{
    "0002/astro-ph0002382_arXiv.txt": {
        "abstract": "We report on $V$ and $R$ high speed photometry of the dwarf nova EX~Dra in quiescence and in outburst. The analysis of the outburst lightcurves indicates that the outbursts do not start in the outer disc regions. The disc expands during the rise to maximum and shrinks during decline and along the following quiescent period. The decrease in brightness at the later stages of the outburst is due to the fading of the light from the inner disc regions. At the end of two outbursts the system was seen to go through a phase of lower brightness, characterized by an out-of-eclipse level $\\simeq 15$ per cent lower than the typical quiescent level and by the fairly symmetric eclipse of a compact source at disc centre with little evidence of a bright spot at disc rim.  New eclipse timings were measured from the lightcurves taken in quiescence and a revised ephemeris was derived. The residuals with respect to the linear ephemeris are well described by a sinusoid of amplitude 1.2 minutes and period $\\simeq 4$ years and are possibly related to a solar-like magnetic activity cycle in the secondary star. Eclipse phases of the compact central source and of the bright spot were used to derive the geometry of the binary. By constraining the gas stream trajectory to pass through the observed position of the bright spot we find $q=0.72\\pm 0.06$ and $i= 85^{+3}_{-2}$ degrees. The binary parameters were estimated by combining the measured mass ratio with the assumption that the secondary star obeys an empirical main sequence mass-radius relation. We find $M_1 = 0.75\\pm 0.15 \\; M_\\odot$ and $M_2 = 0.54\\pm 0.10 \\; M_\\odot$. The results indicate that the white dwarf at disc centre is surrounded by an extended and variable atmosphere or boundary layer of at least 3 times its radius and a temperature of $T\\simeq 28000 \\;K$. The fluxes at mid-eclipse yield an upper limit to the contribution of the secondary star and lead to a lower limit photometric parallax distance of $D= 290 \\pm 80\\; pc$. The fluxes of the secondary star are well matched by those of a M$0\\pm2$ main sequence star. ",
        "introduction": "Dwarf novae are mass-exchanging binaries in which a late type star (the secondary) overfills its Roche lobe and transfers matter to a companion white dwarf (the primary) via an accretion disc. These systems show recurrent outbursts of 2--5 magnitudes on timescales of a few weeks to months caused either by an instability in the mass transfer from the secondary star or by a thermal instability in the accretion disc which switches the disc from a low to a high-viscosity regime (Warner 1995 and references therein). During outburst most of the light arises from the bright, optically thick accretion disc, while in quiescence the dominant sources of light are the white dwarf and the bright spot formed by the impact of the infalling gas stream with the edge of the disc.  Eclipsing dwarf novae are probably the best sites for the study of accretion physics as the occultation of the accretion disc and white dwarf by the secondary can be used to constrain the geometry and parameters of the binary, and tomographic techniques such as eclipse mapping (Horne 1985) and Doppler tomography (Marsh \\& Horne 1988) can be applied to probe the structure and dynamics of the accretion flow.  EX Draconis (= HS1804+67) was detected in the Hamburger Quasar Survey (Bade et~al. 1989) and shown to be an eclipsing dwarf nova with an orbital period of 5.04 hr by Barwig et~al. (1993). From spectroscopic observations made in quiescence, Billington, Marsh \\& Dhillon (1996) found that the secondary star is of spectral type M1 to M2 and that it contributes almost all of the light at mid-eclipse. Their analysis showed that the inner face of the secondary is significantly irradiated by the white dwarf. They found a rotational broadening of $v\\sin i = 140\\; km\\,s^{-1}$ and a radial velocity semi-amplitude of $K_2 = 210\\; km\\,s^{-1}$ for the secondary star which leads to a spectroscopic mass ratio of $q=0.8$ when combined with the $K_1= 167\\; km\\,s^{-1}$ of Barwig et~al. (1993). A relatively small value for the radius of the accretion disc ($0.4\\;R_{L1}$) is derived but no explanation is given of how this estimate was made.  In a follow up study using spectroscopy and photometry of EX~Dra in quiescence and in outburst, Fiedler, Barwig \\& Mantel (1997) measured radial velocity semi-amplitudes of $K_1= 167\\; km\\,s^{-1}$ and $K_2= 223\\; km\\,s^{-1}$ and derived a spectroscopic model for the binary with $q=0.75$, $i= 84.2\\degr$, $M_1= 0.75\\;M_\\odot$ and $M_2= 0.56\\;M_\\odot$. However, the radial velocity curve of the H$\\alpha$ line shows a large phase shift ($\\simeq 0.2$ cycle) with respect to photometric conjunction which casts doubt on the derived value of $K_1$. They use the eclipse phases of the bright spot and white dwarf to derive a photometric mass ratio between 0.7 and 0.8, supporting the spectroscopic model. From the ratios of Ca\\,I and TiO absorption features they infer a spectral type of M0 for the secondary star. Smith \\& Dhillon (1998) use the values of $v\\sin i$ and $K_2$ of Billington et~al. (1996) and the eclipse phase width $\\Delta\\phi$ of Fiedler et~al. (1997) to infer a $K_1= 176 \\;km\\,s^{-1}$.  In this paper we present and discuss high-speed photometry of EX~Dra in quiescence and in outburst. Section~\\ref{obs} describes the observations. In section~\\ref{results} we present and discuss the eclipse lightcurves, provide an updated ephemeris, derive the binary parameters from the eclipse phases of the white dwarf and bright spot, and obtain estimates of the distance to the binary. The results are summarized in section~\\ref{final}. ",
        "conclusions": "\\label{final}  The results of the analysis of $V$ and $R$ high speed photometry of EX~Dra in quiescence and through outburst can be summarized as follows:  \\begin{enumerate}  \\item During the period of the observations EX~Dra showed outbursts with typical amplitudes of $\\simeq 2.0$ mag, duration of $\\simeq 10$ days, and average time between outbursts of $20\\pm 3$ days. The observed amplitudes are larger than those found by Billington et~al. (1996).  \\item The lightcurves during outburst were grouped by outburst phase. The analysis of these lightcurves indicates that the outbursts do not start in the outer disc regions and, therefore, favours the disc instability model. The disc expands during the rise to maximum (as indicated by the increasing width of the eclipse) and shrinks during decline. The decrease in brightness at the later stages of the outburst is due to the fading of the light from the inner disc regions.  \\item At the end of two outbursts the system was seen to go through a phase of lower brightness (named the low state), characterized by the fairly symmetric eclipse of a compact source at disc centre with little evidence of a bright spot at disc rim, and by an out-of-eclipse level $\\simeq 15$ per cent lower than the typical quiescent level.  \\item New eclipse timings were measured from the lightcurves in quiescence and a revised ephemeris was derived. The residuals with respect to the linear ephemeris show a clear cyclical behaviour and can be well described by a sinusoid of amplitude 1.2 minutes and period $\\simeq 4$ years. This period variation is possibly related to a solar-like magnetic activity cycle in the secondary star.  \\item Eclipse phases of the compact central source and of the bright spot were used to derive the geometry of the binary. By constraining the gas stream trajectory to pass through the observed position of the bright spot we find $q=0.72\\pm 0.06$ and $i= 85^{+3}_{-2}$ degrees.  \\item The binary parameters were estimated by combining the measured mass ratio with the assumption that the secondary star in EX~Dra obeys the empirical main sequence mass-radius relation of Smith \\& Dhillon (1998). The set of derived parameters in listed in Table~\\ref{tab5}.  \\item The observed changes in the position of the bright spot with time suggest that the accretion disc shrinks during quiescence by at least $\\simeq 12$ per cent.  \\item The phase of hump maximum is distinct from the radial direction of the bright spot. If the hump maximum is normal to the plane of the shock at the bright spot site then the shock lies in a direction between the stream trajectory and the edge of the disc, making an angle of $41\\degr\\pm 4\\degr$ with the latter.  \\item The white dwarf seems surrounded by an extended, variable atmosphere or boundary layer of at least 3 times its radius. From the derived parameters, a duration of the ingress/egress of the white dwarf of $\\Delta_{wd}= 0.0024\\pm 0.0006$ cycle is predicted.  \\item The fluxes at mid-eclipse of the lightcurves of the low state yield an upper limit to the contribution of the secondary star and lead to a lower limit photometric parallax distance of $D_{MS}= 290\\pm 80\\;pc$.  \\item The fluxes of the central source are well fitted by a white dwarf atmosphere model with $T_{cs}= (28\\pm 3)\\times 10^3\\;K ,\\: \\log g = 8$ and solid angle $\\theta^2_{cs}= [(R_{cs}/R_\\odot)/(D/kpc)]^2= (3.3\\pm 0.5) \\times 10^{-3}$. For a spherical central source, this leads to a distance of $D= 640\\pm 50\\;pc$ if the inner disc is optically thin. The distance estimates from the mid-eclipse fluxes and from the fluxes of the central source can be reconciled if the central source has a toroidal shape with an equatorial diameter of $2\\times 0.037\\;R_\\odot$ and a vertical thickness of $0.012\\; R_\\odot$ (if the inner disc is optically thin) or $0.024\\;R_\\odot$ (if the inner disc is opaque).  \\end{enumerate}  The analysis of the set of lightcurves through outburst with eclipse mapping techniques yields an uneven opportunity to investigate the changes in the structure of an outbursting accretion disc and is the subject of another paper (Baptista \\& Catal\\'an 1999, 2000)."
    },
    "0002/astro-ph0002457_arXiv.txt": {
        "abstract": "We present predictions for the reionization of the intergalactic medium (IGM) by stars in high-redshift galaxies, based on a semi-analytic model of galaxy formation. We calculate ionizing luminosities of galaxies, including the effects of absorption by interstellar gas and dust on the escape fraction $f_{\\rm esc}$, and follow the propagation of the ionization fronts around each galaxy in order to calculate the filling factor of ionized hydrogen in the IGM. For a $\\Lambda{\\rm CDM}$ cosmology, with parameters of the galaxy formation model chosen to match observations of present-day galaxies, and a physical calculation of the escape fraction, we find that the hydrogen in the IGM will be reionized at redshift $z=6.1$ if the IGM has uniform density, but only by $z=4.5$ if the IGM is clumped.  If instead we assume a constant escape fraction of 20\\% for all galaxies, then we find reionization at $z=9.0$ and $z=7.8$ for the same two assumptions about IGM clumping. We combine our semi-analytic model with an N-body simulation of the distribution of dark matter in the universe in order to calculate the evolution of the spatial and velocity distribution of the ionized gas in the IGM, and use this to calculate the secondary temperature anisotropies induced in the cosmic microwave background (CMB) by scattering off free electrons. The models predict a spectrum of secondary anisotropies covering a broad range of angular scales, with fractional temperature fluctuations $\\sim 10^{-7}-10^{-6}$ on arcminute scales. The amplitude depends strongly on the total baryon density, and less sensitively on the escape fraction $f_{\\rm esc}$. The amplitude also depends somewhat on the geometry of reionization, with models in which the regions of highest gas density are reionized first giving larger CMB fluctuations than the case where galaxies ionize surrounding spherical regions, and models where low density regions reionize first giving the smallest fluctuations. Measurement of these anisotropies can therefore put important constraints on the reionization process, in particular, the redshift evolution of the filling factor, and should be a primary objective of a next generation submillimeter telescope such as the Atacama Large Millimeter Array. ",
        "introduction": "The Gunn-Peterson (GP) effect \\cite{gunn65} strongly indicates that the smoothly distributed hydrogen in the intergalactic medium (IGM) is already highly ionized by $z=5$ \\cite{schneider91,lanzetta95}. Barring the possibility of collisional reionization (e.g. Giroux \\& Shapiro 1994), the GP effect implies the presence of very luminous ionizing sources at high redshifts capable of producing enough Lyman continuum (Lyc) photons to cause photoionization of hydrogen by $z\\gsim 5$. The two possible sources of these ionizing photons are QSOs and high mass stars.  Models in which QSOs dominate the production of ionizing photons may be able to meet the GP constraint \\cite{miralda90}. However, such models are strongly constrained by the observed drop in the abundance of bright QSOs above $z\\approx 3$ \\cite{hartwick90,warren94,kennefick95,schmidt95}. Furthermore, \\scite{mhr98} note that a model in which faint QSOs provide all the required ionizing luminosity can be ruled out on the basis of the number of faint QSOs seen in the HDF.  There is, however, growing evidence for the presence of bright galaxies at redshifts as high as $z \\sim 5$ \\cite{spinrad98}, and perhaps even higher \\cite{yahil98}. Thus the other natural candidate sources of ionizing photons are young, high mass stars forming in galaxies at redshifts greater than 5 (e.g. \\pcite{CR86,haiman96,ciardi99}). \\scite{mhr98} note that at $z\\approx 3$ stars in Lyman-break galaxies will emit more ionizing photons into the IGM than QSOs if more than 30\\% of such photons can escape from their host galaxy. Whilst such high escape fractions may not be realistic (e.g. local starbursts show escape fractions of only a few percent, \\pcite{leitherer95}), this does demonstrate that high-redshift galaxies could provide a significant contribution to (or perhaps even dominate) the production of ionizing photons. In this work we will restrict our attention to ionizing photons produced by stars, deferring consideration of the QSO contribution to a later paper.  According to the hierarchical structure formation scenario (e.g. \\pcite{peebles80}) perturbations in the gravitationally dominant and dissipationless dark matter grow, by gravitational instability, into virialised clumps, or halos.  Galaxies, and later stars, then form by the cooling and condensation of gas inside these halos (e.g. \\pcite{wr78,wf91}). Dark matter halos continually grow by merging with other halos (e.g. \\pcite{bower91,bcek}). In the context of this hierarchical scenario, we present a realistic scheme for studying the reionization of the universe by ionizing photons emitted from massive stars.  We focus on the photoionization of the hydrogen component of the IGM. To predict the time dependent luminosity in Lyc photons we use a semi-analytic model of galaxy formation (e.g. \\pcite{kauff93,coleetal94,somerville98}). In particular, we use the semi-analytic model of \\scite{coleetal99}, modified to take into account Compton cooling by cosmic microwave background (CMB) photons, to model the properties of galaxies living in dark matter halos spanning a wide range of masses.  We then estimate the fraction of the ionizing photons which manage to escape each galaxy, and therefore contribute to the photoionization of the intergalactic \\HI. The fraction of ionizing photons escaping is determined on a galaxy-by-galaxy basis, using physically motivated models. Assuming spherical symmetry, we follow the propagation of the ionization front around each halo to compute the filling factor of intergalactic \\HII\\ regions, including the effects of clumping in the IGM.  Finally, using several alternative models for the spatial distribution of ionized regions within a high resolution N-body simulation of the dark matter distribution, we estimate the anisotropies imprinted on the CMB by the patchy reionization process, due to the correlations in the ionized gas distribution and velocities \\cite{sz80,vishniac87}. In previous models many simplifications were made in computing both the spatial distribution of ionized regions and the two-point correlations of gas density and velocity in those regions \\cite{aghanim96,jaffe98,gruzinov98,knox98,peebles98,haiman99}. Our calculations represent a significant improvement over these models as we are able to calculate the two-point correlations between gas density and velocity in ionized regions directly from an N-body simulation.  The rest of this paper is arranged as follows. In \\S\\ref{sec:modified} we outline the features of the semi-analytic model relevant to galaxy formation at high redshifts. In \\S\\ref{sec:fescape} we describe how we calculate the fraction of ionizing photons escaping from galaxies, and observational constraints on the ionizing luminosities and escape fraction at low and high redshift from ${\\mathrm H}\\alpha$ luminosities and \\HI\\ masses and column densities.  In \\S\\ref{sec:fronts} we describe how we calculate the filling factor of photoionized gas in the IGM, including the effects of clumping of this gas. We then present our predictions for reionization, including the effects on the reionization redshift of using different assumptions about escape fractions and clumping factors.In \\S\\ref{sec:variants} we examine the robustness of our results to changes in the other parameters of the semi-analytic galaxy formation model. In \\S\\ref{sec:spatial} we describe how the semi-analytic models are combined with N-body simulations to calculate the spatial distribution of the photoionized IGM. We then calculate the spectrum of anisotropies introduced into the CMB by this ionized gas.  Finally, in \\S\\ref{sec:conc} we summarize our results and examine their consequences. ",
        "conclusions": "\\label{sec:conc}  We have outlined an approach to studying the reionization of the universe by the radiation from stars in high redshift galaxies. We have focussed on the reionization of hydrogen, but the approach can be generalised to study helium reionization (e.g. \\pcite{giroux94}), and also to include radiation from quasars. Our main conclusions are:  (i) Using a model of galaxy formation constrained by several observations of the local galaxy population, enough ionizing photons are produced to reionize the universe by $z=11.7$. This assumes that all ionizing photons escape from the galaxies they originate in, and that the density of the IGM is uniform. Reionization is delayed until $z\\approx 10.9$ in the case of a clumped IGM, in which gas falls into halos with virial temperatures exceeding $10^4$K. Galaxies can reionize such a clumped IGM by $z=5$ providing that, on average, at least 4\\% of ionizing photons can escape from the galaxies where they are produced. In the case of a uniform IGM, an escape fraction of only 1.4\\% is sufficient to reionize by $z=5$.  Using a physical model for the escape of ionizing radiation from galaxies, in which photons escape through ``\\HII\\ chimneys'' ionized in the gas layers in galaxy disks \\cite[hereafter DS94]{doveshull94}, we predict reionization by $z=6.1$ for a uniform IGM or by $z=4.5$ for a clumped IGM. Models which assume that all the gas in galaxy disks remains neutral are unable to reionize even a uniform IGM by $z=0$. Using alternative estimates of the IGM clumping factor from \\scite{gnedin97} or \\scite{valageas99}, we find reionization redshifts comparable with those found using our own clumping model, i.e. in the range $z=$ 4.5--5.0 with the DS94 model for the escape fraction.  (ii) Once the ionizing escape fraction and IGM clumping factor have been specified, our estimates for the filling factor of ionized gas in the IGM are reasonably robust, providing that we consider only models which are successful in matching the H$\\alpha$ luminosity function of galaxies at $z=0$. By far the greatest remaining influences on the ionized filling factor come from the value of the baryon fraction $\\Omega_{\\rm b}$ and the prescription for feedback from supernovae. However, we have shown that altering these parameters also produces large changes in the $z=0$ H$\\alpha$ luminosity function.  (iii) We combined our model for reionization with N-body simulations of the dark matter distribution in order to predict the spectrum of secondary anisotropies imprinted on the CMB by the process of reionization. The {\\it shape} of this spectrum is almost independent of the assumptions about reionization, but the {\\it amplitude} depends on the spatial distribution of the ionized regions, the redshift at which reionization occurs and the baryon fraction. We find considerably more power in the anisotropy spectrum at small $\\ell$ than predicted by models which do not account for the large-scale correlations in the gas density and velocity produced by gravity. Despite the uncertainty in the spatial distribution of ionized regions, we are able to determine the amplitude of this spectrum to within a factor of three for a given $\\Omega_{\\rm b}$ (the amplitude being proportional to $\\Omega_{\\rm b}^2$). The results found by \\scite{bruscoli00} using a similar technique are reasonably consistent with ours, once differences in $\\Omega_{\\rm b}$ and other cosmological parameters are allowed for.  Detection of these secondary anisotropies, which would constrain the reionization history of the Universe, would require fractional temperature fluctuations of $\\sim 10^{-7}$ to be measured on angular scales smaller than several arcminutes. Although the Planck and MAP space missions are unlikely to have sufficient sensitivity to observe such anisotropies, the Atacama Large Millimeter Array is expected to be able to measure temperature fluctuations of the level predicted at $\\ell \\sim 10^4$ in a ten hour integration.  Previous studies of reionization have either used an approach similar to our own, i.e. employing some type of analytical or semi-analytical model (e.g. \\pcite{haiman96,valageas99,chiu99,ciardi99}), or else have used direct hydrodynamical simulations (e.g. \\pcite{gnedin97}). While the latter technique can in principle follow the detailed processes of galaxy formation, gas dynamics and radiative transfer, in practice the resolutions attainable at present do not allow such simulations to resolve the small scales relevant to this problem. Furthermore, the implementation of star formation and feedback in such models is far from straightforward.  There are two main uncertainties in our approach, as in most others: the fraction \\fesctxt\\ of ionizing photons that escape from galaxies, and the clumping factor $f_{\\rm clump}$ of gas in the IGM.  Future progress depends on improving estimates of the effects of clumping using larger gas dynamical simulations, on better modelling of the escape of ionizing photons from galaxies, and on better understanding of star formation and supernova feedback in high redshift objects."
    },
    "0002/astro-ph0002511_arXiv.txt": {
        "abstract": "\\begin{small} We analyze high resolution, high signal-to-noise spectra of six red-giant-branch (RGB) stars in the globular cluster M~3 (NGC~5272) and three in M~13 (NGC~6205) that were obtained with the Mayall 4-meter telescope and echelle spectrometer on Kitt Peak. The spectra include lines of O, Na, Mg, Al, Si, Ca, Ti, V, Mn, Fe and Ni. We also analyze the [Al/Fe] values of 96 RGB stars in M~13 covering the brightest 3.5 magnitudes, which include 66 measurements that were derived from moderate resolution, low signal-to-noise spectra obtained with the WIYN 3.5-meter telescope and Hydra multi-object spectrograph, also on Kitt Peak. In addition, we compile from the literature and inspect the [Na/Fe] values of 119 RGB stars in M~13. We test for bimodality in the [Al/Fe] and [Na/Fe] distributions using the KMM algorithm and find that the [Al/Fe] values in M~13 are distributed bimodally at all points along the RGB that were observed, while the [Na/Fe] values are bimodal only over the brightest two magnitudes. The ratios of Al-enhanced to Al-normal and Na-enhanced to Na-normal giants increase towards the tip of the RGB in M~13, which is suggestive of deep mixing in this cluster. The limited M~3 data exhibit a bimodal distribution of [Al/Fe] values and are suggestive of no deep mixing; however, they are too few to be conclusive. We further test for a relationship between deep mixing on the RGB and a second parameter that can create the extended blue tail seen along the  horizontal-branches of some clusters by using an ``instantaneous'' mixing algorithm, which we develop here. We conclude that the data for both clusters are consistent with deep mixing as a ``blue-tail second parameter'', and we suggest future observations to further constrain the results. Finally, we offer a solution to the problem of over producing sodium during deep mixing that is based on the depletion of $^{22}$Ne in asymptotic-giant-branch stars and suggest that pollution might best be traced by {\\spr} elements in the Sr-Y-Zr peak. \\end{small} ",
        "introduction": "According to canonical stellar evolution models, the by-products of the nuclear processing around the hydrogen-burning shell (H~shell) of low-mass red-giant-branch (RGB) stars should remain confined to the stellar interior; however, observations over the past 25 years have shown star-to-star variations in the elements C, N, O, Na, Mg and Al, among others, on the surfaces of globular cluster red giants (see Kraft 1994, Briley et al. 1994 and Cavallo 1998a for detailed reviews of the observations). In particular, the data show evidence of the CNO cycle that dominates the energy production in such stars: C and O are anticorrelated with N, while the {\\cratio} is near the equilibrium value of 4 in many clusters \\citep{SS91,S96b,BSKL97,BSSBN97,ZWB96}. While the first dredge-up phenomenon \\citep{Iben67} does alter the carbon and nitrogen abundances slightly, it cannot account for the observed large variations of these elements and their isotopic ratios, nor can it account for the variations of the other elements. In addition, some elements show evidence for gradual changes along the RGB, indicating that something is occurring during the course of evolution to facilitate these alterations. For example, C becomes more depleted with decreasing $V$ in the clusters M~15, M~55, M~92 and NGC~6397 \\citep{BDG79,Carbon1982,TCLSK83, BBHD90}.  Two separate approaches have been developed to address the observations. One assumes that some form of non-canonical mixing occurs along the RGB, which gradually brings material from around the H~shell to the stellar surface \\citep[hereafter, SM79]{SM79}. Models by SM79, \\citet{DD90}, \\citet{LHS93}, \\citet{CSB96}, \\citet{DW96} and \\citet[hereafter, CSB98]{CSB98} have shown that most variations along the brighter part of the RGB can be explained by nuclear processing around the H~shell combined with mixing. The source of mixing is generally assumed to be rotationally induced meridional circulation currents (SM79); although, other theories abound \\citep{LHZ97,FAK99}. The observations by \\citet{Peterson83} that show the horizontal-branch (HB) stars in M~13, a cluster with large variations of oxygen and aluminum on the RGB, rotating nearly a factor of two faster than the HB stars in M~3, a cluster with a composition similar to M~13, but with less extreme abundance variations along it's RGB, support the SM79 hypothesis.  The second approach assumes that some of the variations, particularly those of the heavier elements, are primordial in nature, perhaps originating in the processed envelopes of intermediate-mass asymptotic-giant-branch (AGB) stars that were shed into the nascent cluster environment \\citep{CD81}. While it has been shown that this scenario cannot account for all the variations \\citep{DWW97}, some aspects of it are plausible in light of the data. For example, observations of CN-band strength and sodium variations on the upper main sequence of 47~Tuc, sodium enhancements on the subgiant branch of M~92 and enhancements in the neutron-capture elements in some clusters all point to primordial origins \\citep{BBSH89,BHB91,S96a,BSSLBH96,CCBHS98,KSB98,Ivans99}. The most likely solution to the abundance anomaly problem probably involves a combination of both scenarios, where primordial pollution is present in the cluster, but mixing later plays a role in adjusting the abundance patterns (see, e.g., Denissenkov et al. 1998 and Briley et al. 1999), an approach we examine here.  This paper focuses on determining the chemical abundances in the red giants of the globular clusters M~3 (NGC~5272) and M~13 (NGC~6205) from high resolution, high signal-to-noise echelle spectra obtained with the Mayall 4-meter telescope on Kitt Peak. We choose these two clusters because they are often considered a classical ``second-parameter'' pair since they have markedly different HB's, despite having similar [Fe/H]\\footnote{ We use the usual notation: [X/Y]~=log(X/Y)$_{\\star}$~$-$~log(X/Y)$_{\\sun}$.} values, the first parameter. We discuss the hypothesis that deep mixing along the RGB, which we define as mixing that penetrates the H~shell, brings helium to the surface and affects the HB morphology as a second parameter that creates the extended blue tail in M~13 \\citep{AVS97a,AVS97b}. One oft-quoted choice for the second parameter is a relative age difference between M~3 and M~13 \\citep{FPB97,SCD97,Chaboyer98}; however, this is not borne out by the photometry \\citep{JB98,vdB99,FG99} and leaves open the need for a qualified alternative.  While M~13 is by far the most well-studied cluster for abundance variations, the data for M~3 are lacking. One goal of this paper is to increase our knowledge of the chemical abundances in this latter cluster so that it can be compared with M~13 in greater detail.  The outline of this paper is as follows: we describe the observations, data reduction approach, abundance analysis technique and abundance results in sections 1, 2, 3 and 4, respectively. In section 5 we discuss evidence for and the implications of deep mixing along the RGB, and we give our conclusions in section 6. In the appendix we derive the instantaneous mixing algorithm that is used in section 5. ",
        "conclusions": "Before we discuss our final conclusions, we first remind the reader of the number of assumptions that have gone into our analysis. First, there are errors associated with the abundance determinations that we tried to characterize by allowing for significant variations in the model atmosphere parameters, which contribute the most to the uncertainty in the analysis. Second, the inclusion of the WIYN data into our analysis comes at a price: the data have poor signal-to-noise ratios, come from only one line and require the assumption that, for some stars with indeterminate line strengths, the [Al/Fe] value is ``low.'' Third, despite that fact that this is the largest compilation of [Al/Fe] and [Na/Fe] values in one globular cluster to be analyzed in a single paper, the data are still subject to small number effects, particularly at the RGB tip. Unfortunately, there are only so many tip stars that can be spectroscopically measured from the ground, leaving this problem difficult to solve. We suggest the best way to handle the small numbers is to expand this analysis to other clusters for a broad comparative study. Fourth, the models have many assumptions in them: we assume that canonical evolution holds and add in our mixing algorithm after the fact, we assume that mixing is instantaneous, we assume that the abundances are distributed as per \\citet{DDNW98}, we assume that our reaction rates are accurate, and we assume that we adequately searched the parameter space allowed by the uncertainties in the initial abundances, nuclear reaction rates and mass-loss rates. Fifth, we assume that no other second parameter affects the relationship between the M~3 and M~13 HB morphologies. Sixth, we make no attempt to correct for blending of the AGB with the RGB when performing our analysis. Approximately 20\\% of the red giants above the point where giant branches merge are supposedly AGB interlopers based on comparative lifetimes: the problem is to determine which ones are really AGB stars. This might not be as much of a problem for the M~13 sample, however, since blue HB stars tend to evolve away from the AGB. The best workaround for this problem is also an extension of our analysis to other clusters to look for consistent trends despite this, and the other, uncertainties.  The importance of having accurate nuclear reaction rates cannot be overstated. This is particularly true when using aluminum as a diagnostic of deep mixing. If we were to vary, for example, the $^{26}$Mg proton-capture rate to its upper limit in range of $T_9~=~0.05-0.06$, the production of Al can move outside the H shell, although, just barely. Depending on the initial abundances of $^{25}$Mg and $^{26}$Mg, this might be able to account for the full enhancements of aluminum that we observed. In addition, according to the NACRE compilation, the rate for $^{26}$Al$^{g}(p,{\\gamma})^{27}$Si is uncertain by as much as three orders of magnitude in the same temperature range. Increasing these rates might help solve the problem presented by \\citet{CDW99} who show that, if mixing occurs below the top of the H shell, sodium is overproduced due to the extra enhancement from $^{20}$Ne in the NeNa cycle, a result we confirm with our instantaneous mixing algorithm. If the $^{26}$Mg proton-capture rate is near its upper limit, then deep mixing is not required to produce the observed aluminum abundances and sodium is not over enhanced compared to the observations.  Our general results for the M~3 and M~13 abundances obtained in this work show the usual trends in the proton-capture, ${\\alpha}$ and iron-peak elements: the sodium and aluminum abundances are anticorrelated with oxygen, the ${\\alpha}$ elements are enhanced by approximately 0.3 dex and the iron-peak elements remain constant.  Our analysis shows that the variation in both the [Al/Fe] and [Na/Fe] ratios are consistent with deep mixing occurring on the RGB in M~13 and not in M~3. The aluminum and sodium data are correlated for the M~13 giants; although, the Al ratio is probably a better indicator of deep mixing since it is more easily separated into ``high'' and ``low'' groups. We would not expect such a similar tight correlation between aluminum and sodium in the M~3 giants since sodium can be enhanced without increasing the aluminum abundance if the mixing currents do not penetrate the H shell, as seems indicated in M~3 from the low Al ratio. However, some semblance of a correlation between aluminum and sodium might be set up by primordial effects in this cluster. In addition, the Na ratio increases near the same magnitude as the Al ratio, which is contrary to the previous predictions that sodium should be enhanced further down the RGB from $^{22}$Ne (CSB98). Our models show that this would be expected if the $^{22}$Ne were depleted in primordial intermediate-mass AGB stars.  When comparing the Al ratio with the HB ratio, it seems that the assumption of deep mixing as a blue-tail parameter is self-consistent; however, the large range allowed in the actual number of mixed RGB stars and the empirical definitions of ``blue'' and ``red'' HB stars do not constrain the results enough to be firmly conclusive. Again we suggest that a similar analysis as the one presented here be extended to other clusters to determine the Al ratio as a function of $V$ and to compare this with the HB ratio. If the Al ratio at the RGB tip can be shown to be a predictor of the HB ratio, then helium mixing would certainly be given greater credence as a blue-tail second parameter, supplanting the oft-assumed cluster age differences that have been shown to fail for this classical pair of clusters. In particular, we suggest the study of metal-rich clusters to see if the aluminum distribution is bimodal, and if it is, if the Al ratio varies. According to our models, it should not vary since aluminum cannot be produced in metal-rich cluster giants on the same scale as it can in the intermediate metallicity and metal-poor giants. In addition, we suggest further examination of the sodium abundance in clusters to search for similar behavior as in M~13. Also, we suggest a more extensive comparison of the {\\spr} abundances with the aluminum data as a test of primordial contamination.  Finally, we conclude that the problem of abundance anomalies in globular cluster red giants requires detailed study of the abundance yields from primordial AGB stars as well as an in-depth and complete study of the hydrodynamical evolution of rotating RGB stars. In the meantime, aluminum, and to a lesser extent, sodium, give the best diagnostics of deep mixing during the evolution up the RGB and the {\\spr} elements near the Sr-Y-Zr peak are the best tracers of AGB pollution from IMS."
    },
    "0002/astro-ph0002277_arXiv.txt": {
        "abstract": "We present sensitive  sub-mm imaging of the Seyfert  1 galaxy NGC 7469 at 850  $\\mu $m and 450  $\\mu $m with the  Submillimetre Common User Bolometer Array (SCUBA) on  the James Clerk Maxwell Telescope (JCMT) and $ ^{12}$CO J=3--2  line observations of its central starbursting region.  The global dust spectrum,  as constrained by the new set of sub-mm data and  available 1.30~mm and IRAS 100~$\\mu  $m, 60 $\\mu $m data reveals  a dominant warm  dust component with a  temperature of $\\rm T_{\\rm d}  \\sim 35$ K and a  global molecular gas-to-dust ratio $\\rm M(H_2)/M_{\\rm d}\\sim 600$.   Including the atomic gas component yields a  total gas-to-dust ratio  of $\\sim 830$.  Such  high values are typical  for IR-bright  spirals and in  order to  reconcile them with the  significantly lower ratio  of $\\sim 100$ obtained  for the Milky Way  a cold dust reservoir, inconspicuous  at FIR wavelengths, is usually  postulated.  However, while  there is good  evidence for the presence of cold gas/dust in NGC~7469 beyond its central region, our 450 $\\mu $m map  and available interferometric $ ^{12}$CO J=1--0 maps show the bright sub-mm/CO  emission confined in the inner $\\sim 2.5$~kpc,  where a  high $  ^{12}$CO (J=3--2)/(J=1--0)  ratio ($\\sim 0.85-1.0$) is  measured.  This is  consistent with molecular  gas at $\\rm T_{\\rm kin}\\ga 30$K, suggesting that the bulk of the ISM in the starburst   center   of   NGC~7469   is  warm.    Nevertheless   the corresponding  total gas-to-dust  ratio there  remains high,  of the order of  $\\sim 500$.  We  argue that, rather than  unaccounted cold dust mass, this high ratio  suggests an overestimate of $\\rm M(H_2)$ from its associated $ ^{12}$CO J=1--0 line luminosity by a factor of $\\sim  5$  when a  Milky  Way value  for  this  conversion is  used. Finally the diffuse  cold gas and dust that is  the likely source of the  observed faint  extended  450  $\\mu $m  and  $ ^{12}$CO  J=1--0 emission has an estimated  total gas-to-dust ratio of $\\sim 50-160$, closer to the Galactic value. ",
        "introduction": "The role of  molecular gas as the fuel of  both starburst activity and an Active Galactic Nucleus (AGN)  is now well established.  In Seyfert galaxies such gas is found on scales ranging from the inner $\\sim 1-2$ kpc ``feeding''  an intense circumnuclear starburst down  to $\\rm L\\la $100~pc from the AGN, where in the form of a geometrically thick torus obscures the active nucleus along certain viewing angles thus creating the difference between type 2 and type 1 Seyferts (Miller \\& Antonucci 1983; Krolik 1990 and references therein).   Estimates  of  molecular  gas   mass  in  other  galaxies  from  their velocity-integrated $ ^{12}$CO J=1--0 luminosity $\\rm L_{\\rm CO}$ (in K km s$ ^{-1}$ pc$ ^2$)  are based on the so-called conversion factor $\\rm  X_{\\rm  CO}=M(H_2)/L_{\\rm  CO}\\approx  5\\ M_{\\odot  }  (K\\  km\\ s^{-1}\\  pc^2)^{-1}$ (e.g. Solomon  \\& Barrett  1991) whose  value is derived from studies of molecular gas in the Galactic disk.  Numerous studies quantify  the effects of  the various physical  conditions on this   factor  (e.g.    Bryant  \\&   Scoville  1996;   Israel  1997). Particularly  in the starburst  nuclei of  very luminous  IR galaxies ($\\rm  L_{\\rm  FIR} >  10^{11}\\  L_{\\odot  }$)  these conditions  are significantly  different than  in the  disk environment  of quiescent spirals like  the Milky  Way.  More specifically  in the  inner $\\sim 1-2$ kpc,  where a starburst  usually occurs, the  gas differentiates into  two distinct phases  (e.g.  Aalto  et al.   1995) with  a warm, diffuse  and possibly  non  self-gravitating phase  dominating the  $ ^{12}$CO  emission.  This  results in  a significant  overestimate of $\\rm H_2$  mass from $\\rm L_{\\rm  CO}$ when a Galactic  value of $\\rm X_{\\rm CO}$ is used (Solomon et al. 1997; Downes \\& Solomon~1998).    The aforementioned  conditions are expected  in the nuclear  region of the SBa Seyfert~1  galaxy NGC~7469 (Arp 298, Mrk  1514), a luminous IR source  with  $\\rm L_{\\rm  FIR}\\approx  3\\times  10^{11}\\ L_{\\odot  }$ emanating  from a  powerful starburst  deeply embedded  into  its $\\rm L\\sim 1$~kpc circumnuclear region (Wilson  et al.  1991; Genzel et al. 1995).  High resolution $  ^{12}$CO J=1--0 observations revealed large amounts of  molecular gas  within this region  (Meixner et  al.  1990; Tacconi \\&  Genzel 1996)  with mass comparable  to the  dynamical, and there  is  significant  evidence  that  application  of  the  standard Galactic   conversion  factor  overestimates   $\\rm  M(H_2)$   in  the circumnuclear starburst of this galaxy (Genzel et al.  1995).  In this paper we present  sensitive sub-mm imaging of NGC  7469 at 850~$\\mu $m and  450~$\\mu $m  and $  ^{12}$CO J=3--2  spectroscopy of  its central region.   Under the  assumption of  a canonical  gas-to-dust  ratio of $\\sim 100$,  our results  confirm the overestimate  of H$_2$  gas mass when a  Galactic value for $\\rm  X_{\\rm CO}$ is  used, and demonstrate the significance of  CO spectroscopy and sub-mm imaging  in offering a better assessment of  the molecular gas mass and  its average physical conditions.   Throughout  this work we  adopt $\\rm  H_{\\circ }  = 75\\  km\\ sec^{-1}\\ Mpc^{-1}$ and $\\rm  q_{\\circ } = 0.5$, which for  cz=4900 km s$ ^{-1}$ yields a  luminosity distance  of $\\rm D_{\\rm  L}\\sim 66$ Mpc  for NGC 7469, where 1$''$ corresponds to $\\sim 310$ pc. ",
        "conclusions": "We presented sensitive 850 $\\mu $m,  450 $\\mu $m maps of the Seyfert 1 galaxy  NGC 7469  and a  measurement of  the $  ^{12}$CO  J=3--2 line towards its starbursting central region. Our main conclusions can be summarized as follows  1. The FIR/sub-mm/mm  spectrum of this  source is dominated by  a warm dust component  with a  temperature of $\\sim  35$ K and  a global molecular gas-to-dust ratio of $\\sim 600$, both typical for IR-luminous spirals.  Including the atomic gas mass  yields a total gas-to-dust ratio of $\\sim 800$.    2. The warm  dust and gas   lies in the inner   8$''$ ($\\sim $2.5 kpc) where a  starburst is embedded.   In this region  we find no evidence for  a significant  mass of  cold dust, yet   the ratio of the mainly molecular gas to the dust mass is $\\sim 500$, still five times larger than the  Galactic value.  We argue  that this  is   the result  of a systematic overestimate  of $\\rm  H_2$ mass  by a  factor of $\\sim 5$ when a Galactic value for $\\rm X_{\\rm  CO}=M(H_2)/L_{\\rm CO}$ is used in starburst environments.   3. On larger  scales (radius  $\\ga 1.2$  kpc) the ISM  in NGC  7469 is dominated by cold, sub-thermally excited  gas where the faint 450 $\\mu $m and $ ^{12}$CO J=1--0 emission originate. The gas-to-dust ratio for this   phase,  is   $\\sim   50-150$;  in   better   accord  with   the Milky~Way~value.   \\subsection"
    },
    "0002/astro-ph0002041_arXiv.txt": {
        "abstract": "We present a convolution-based algorithm for finding cosmic rays in single well-sampled astronomical images.  The spatial filter used is the point spread function (approximated by a Gaussian) minus a scaled delta function, and cosmic rays are identified by thresholding the filtered image.  This filter searches for features with significant power at spatial frequencies too high for legitimate objects. Noise properties of the filtered image are readily calculated, which allows us to compute the probability of rejecting a pixel not contaminated by a cosmic ray (the false alarm probability).  We demonstrate that the false alarm probability for a pixel containing object flux will never exceed the corresponding probability for a blank sky pixel, provided we choose the convolution kernel appropriately.  This allows confident rejection of cosmic rays superposed on real objects.  Identification of multiple-pixel cosmic ray hits can be enhanced by running the algorithm iteratively, replacing flagged pixels with the background level at each iteration. ",
        "introduction": "Images from most current-day astronomical instruments have tractable noise properties.  An exemplary case is optical images from CCD detectors, whose uncertainties are generally dominated by the Poisson statistics of the detected photons, with (usually smaller) contributions from detector read noise, dark current, and other comparatively minor nuisances.  Most of these noise sources are well approximated by Gaussian distributions, and their sum is therefore also well approximated by a Gaussian.  Cosmic rays impinging on a detector can yield large signals over single pixels or small groups of pixels, thereby introducing a distinctly non-Gaussian tail to the noise distribution.  The most common approach to removing cosmic rays from astronomical images is to take multiple exposures and combine them with some sort of outlier rejection.  Real astronomical objects should (usually) be present on multiple frames, while cosmic ray hits will not generally repeat. Such methods have been presented in the literature by (e.g.) Shaw \\& Horne (1992) and Windhorst, Franklin, \\& Neuschaefer (1994), and are widely implemented in astronomical image processing packages.  However, there are times when multiple images are not available, or when the sources of interest may be moving or varying on timescales short compared to the interval between exposures.  In these cases, a cosmic ray rejection method capable of operating on single exposures is necessary.  Cosmic ray rejection in single frames can also be useful even when multiple exposures are to be stacked, since stacking often requires spatial interpolation of the input images, and any cosmic rays not previously identified can be spread over many pixels by spatially extended interpolation kernels.  Additionally, if a stack of images has widely different point spread function (PSF) widths, rejection algorithms used while stacking tend either to be overly lenient, potentially admitting cosmic rays; or overly strict, discarding valid data from images with very good or very bad seeing. Examples of both these behaviors are offered by sigma clipping algorithms, where the contribution of a particular exposure to a stack is discarded if it differs from the mean (or median) intensity at that location by more than $k \\sigma$, where $k$ is a constant (generally with $2 \\la k \\la 5$) and $\\sigma$ measures the intensity uncertainty at that location.  If $\\sigma$ is measured directly from the list of exposure intensities at a fixed sky position, a lenient rejection results, while if $\\sigma$ is taken from the known Poisson statistics of electrons in single exposures, a strict rejection results.  To identify cosmic rays in single exposures, rejection algorithms rely on the sharpness of cosmic rays relative to true astronomical objects. That is, any legitimate object in our astronomical image is blurred by the PSF, but there is no such requirement on cosmic ray hits. Provided the image is well-sampled (in practice, $\\ga 2$ pixels across the PSF full width at half maximum), cosmic ray hits can be identified as those features with spatial variations too rapid for consistency with the PSF.  Murtagh (1992) and Salzberg et al (1995) have explored trainable classifier approaches to single-image cosmic ray rejection. Their methods have the advantage of applicability to substantially undersampled data (from the WF/PC-I instrument on the Hubble Space Telescope).  On the other hand, these methods ultimately rely on a training set, which may be subjectively defined.  The present paper explores a method suggested by Fischer and Kochanski (1994), who remark that the optimal filter for detecting [single-pixel] cosmic ray hits is the point spread function minus a delta function.  This can be regarded as a difference between the matched filter for detecting point sources (i.e. the PSF) and that for detecting single pixels (i.e. a delta function).  There is one free parameter in such a filter, which is the amplitude ratio of the two functions.  We develop this filtering method in detail by considering the cosmic ray rejection rates and false alarm rates.  Much of our analysis is devoted to choosing the delta function amplitude appropriately.  With a careful choice of this parameter, it is possible to ensure that the false alarm rate nowhere exceeds its value in blank sky regions.  In section~\\ref{math}, we derive the noise properties of our filtered image, and explain how to tune the filter to avoid excessive rejection of valid data.  In section~\\ref{iraf}, we discuss practical issues that arise when implementing our algorithm. Section~\\ref{simulations} presents simulations used to verify the algorithm's performance. Finally, in section~\\ref{theend} we summarize our work, describe our usual application for our algorithm, and comment on a desirable future direction for cosmic ray rejection algorithms. ",
        "conclusions": "\\label{theend} We have presented a cosmic ray rejection algorithm based on a convolution of the input image.  The advantages of the method spring from the linear nature of the spatial filter, which allows us to determine the noise properties of the filtered image and so to calculate and control the probability of rejecting the central pixel (or indeed any pixel) of a point source.  This safety mechanism ensures that cosmic ray rejection can be applied throughout the image, without special treatment for the locations of sources.  The sensitivity to cosmic rays is of course reduced at the locations of objects, because of the added Poisson noise contributed by object photons and the resulting need to maintain a positive expectation value in the filtered image.  We usually apply our method conservatively, considering pixels innocent until proven guilty beyond any reasonable doubt.  This means that given some uncertainty in the measured point spread function, we use a convolution kernel that is slightly narrower than our best estimate of the PSF (generally by about 10\\%).  This choice depends on the relative importance of keeping legitimate sources and rejecting spurious ones for the scientific problem at hand.  Our original goal in developing this algorithm was to flag and replace cosmic ray hits in individual exposures that are later aligned and stacked.  The alignment procedure requires interpolating the original images, and we use sinc interpolation to preserve the spatial resolution and noise properties of the input image.  However, sinc interpolation assumes well sampled data and responds badly to cosmic rays, spreading their effects over many more pixels than were originally affected and motivating us to replace them at an early stage.  We nevertheless have a second chance to reject cosmic rays by looking for consistency among our different exposures when we stack them, and this second chance helps motivate our generally conservative approach to cosmic ray flagging.  By applying the algorithm developed here followed by sigma rejection during image stacking, we exploit two distinct properties of cosmic rays: They are sharper than the point spread function, and they do not repeat from exposure to exposure.  However, we are using these two tests in sequence.  An algorithm exploiting both pieces of information simultaneously could potentially yield more sensitive cosmic ray rejection.  For general data sets, such an algorithm would have to handle stacks of unregistered images with different PSFs, making its development difficult but potentially rewarding.  An interesting effort in this regard is Freudling's (1995) algorithm, which identifies cosmic rays in the course of deconvolving and coadding images with Hook \\& Lucy's (1992) method."
    },
    "0002/astro-ph0002088_arXiv.txt": {
        "abstract": "The properties of old globular cluster systems (GCSs) in galaxy halos offer unique insight into the physical processes that conspire to form any generic star cluster, at any epoch. Presented here is a summary of the information obtained from (1) the specific frequencies (total populations) and spatial structures (density vs.~galactocentric radius) of GCSs in early-type galaxies, as they relate to the efficiency (or probability) of bound cluster formation, and (2) the fundamental role of a scaling between cluster mass and energy among Galactic globulars in setting their other structural correlations, and the possible implications for star formation efficiency as a function of mass in gaseous protoclusters. ",
        "introduction": "Until quite recently, it was commonly assumed that the old globular clusters in galaxy halos were the remnants of a unique sort of star formation that occurred only in a cosmological context. The discovery of young, massive, ``super'' star clusters in local galaxy mergers and starbursts has clearly done much to change this perception; but at least as important is the parallel recognition that star formation in the Milky Way itself proceeds---under much less extreme conditions---largely in a clustered mode. Observations of entire starbursts (Meurer et al.~1995) and individual Galactic molecular clouds (e.g., Lada 1992), as well as a more general comparison of the mass function of molecular cloud clumps and the stellar IMF (Patel \\& Pudritz 1994), all argue convincingly that (by mass) most new stars are born in groups rather than in isolation. The production of a true stellar cluster---one that remains bound even after dispersing the gas from which it formed---is undoubtedly a {\\it rare} event, but it is an exceedingly {\\it regular} one.  Seen in this light, the globular cluster systems (GCSs) found in most galaxies can be used to good effect as probes not only of galaxy formation but also of an important element of the generic star-formation process at any epoch. This is arguably so even in cases where newly formed clusters may not be ``massive'' according to the criteria of this workshop (the main issue being simply the formation of a {\\it self-gravitating} stellar system), and even though GCSs have been subjected to $10^{10}$ yr of dynamical evolution in the tidal fields of their parent galaxies (see O.~Gerhard's contribution to these proceedings, and note that theoretical calculations geared specifically to conditions both in the Milky Way [Gnedin \\& Ostriker 1997] and in the giant elliptical M87 [Murali \\& Weinberg 1997] suggest that GCS properties are most affected by evolution inside roughly a stellar effective radius in each case). ",
        "conclusions": ""
    },
    "0002/astro-ph0002331_arXiv.txt": {
        "abstract": "We report new interferometric observations of the $\\COjtwo$ rotational transition on Titan.  We find that the spectrum is best fit by a uniform profile of 52 ppm, with estimated errors of 6 ppm (40 to 200 km) and 12 ppm (200 to 300 km). ",
        "introduction": "The atmosphere of Titan exhibits a complex photochemistry, and many nitriles and hydrocarbons have been detected by Voyager spacecraft and from Earth.  Until recently, however, only two oxygen-bearing species had been detected on Titan: $\\COtwo$ (observed by Voyager 1; \\cite{samu1983}) and CO (observed from Earth; \\cite{lutz1983}).  The presence of oxygenated molecules is interesting because the atmosphere of Titan is strongly reducing.  The cold temperatures of the lower stratosphere and the troposphere imply that $\\COtwo$ condenses out of the lower atmosphere and is continuously deposited on the surface.  To sustain the carbon dioxide abundance a source of oxygen is needed, and it is generally assumed to be supplied in water from bombardment of the upper atmosphere by icy grains.  In this model vaporized water is quickly photolyzed to produce OH, and OH reacts with hydrocarbon radicals such as $\\CHthree$ to produce CO.  CO in turn reacts with OH to produce $\\COtwo$ (\\cite{samu1983}, \\cite{yung1984}, \\cite{toub1995}, \\cite{lara1996}). While $\\COtwo$ has a short lifetime (order 10$^3$--10$^4$ years), the photochemical lifetime of $\\CO$ in the atmosphere of Titan is estimated to be very long ($\\sim 10^9$ years; \\cite{yung1984}, \\cite{chas1991}).  Observationally the missing piece of the oxygen chemistry has been the source, water.  Recently, water vapor was detected in the upper atmosphere of Titan by the Short Wavelength Spectrometer (SWS) aboard the Infrared Satellite Observatory (\\cite{cous1998}).  With observations of the three major components of oxygen chemistry, it is now possible to check the internal consistency of photochemical models, and to compare the oxygen chemistry and water infall rate of Titan with the other giant planets, particularly Saturn (\\cite{feuc1997}, \\cite{cous1998}).  Understanding the oxygen chemistry relies on accurate knowledge of the abundance and distribution of each species.  A longstanding discussion regarding the CO distribution in Titan's atmosphere, spanning more than a decade, has been primarily directed toward determining if CO is well-mixed (\\cite{mart1988}, \\cite{gurw1995}, \\cite{hida1998}).  Since the residence lifetime of CO is long compared to transport timescales, the molecular weight of $\\CO$ is the same as for the dominant $\\Ntwo$ gas, and the atmosphere is never cold enough for CO to condense, carbon monoxide should be uniformly mixed in the Titan atmosphere to high altitudes.  Observational data, however, give conflicting results.  Table 1 provides data on the CO abundance as measured by ground-based observers over the past 17 years.  These observations have been sensitive to either the troposphere (near- and mid-IR) or the stratosphere (millimeter).  The data in Table 1 show that no clear consensus has emerged regarding the CO abundance, either in the troposphere or the stratosphere.  In this Note, we present an analysis of new interferometric observations of the $\\COjtwo$ line on Titan.  The results of this study have important implications for our understanding of the oxygen budget and photochemistry of the stratosphere of Titan. ",
        "conclusions": "The results presented here are nearly identical to our previous estimate of the CO distribution based on observations of the $\\COjone$ transition (\\cite{gurw1995}) and consistent with the original measurement of tropospheric CO (\\cite{lutz1983}).  Taken together, these measurements suggest a vertical profile of CO that is constant with altitude, at about 52 ppm, from the surface to at least 300 km.  These results are at odds with the recent measurements of Noll (1996), who found a tropospheric abundance of 10 ppm, and Hidayat \\etal (1998), who found a stratospheric CO abundance of around 27 ppm (Table 1).  Noll (1996) explored the possibility that their simple reflecting layer was not the surface, but a higher altitude 'haze' layer.  If the reflecting layer was at 0.9 bar (14 km) the spectrum was best fit with a CO abundance of 60 ppm.  However, based on other evidence they found this model less satisfactory than a surface reflecting layer.  The results of Hidayat \\etal come from an analysis of several lines of CO, including the $\\COjone$ and $\\COjtwo$ lines; the discrepancy between their results and ours does not appear to be due to differences in modeling the atmosphere of Titan, but derives from differences in the measurement techniques and the resulting calibrated spectra (A.~Marten, personal communication).  However, we do point out that the interferometric method does offer advantages over single-dish observations for measuring the very broad lines of CO from the atmosphere of Titan.  We find the model of a uniform distribution of CO in the atmosphere of Titan provides a good fit to our data, but we cannot rule out a difference between the tropospheric and stratospheric CO abundance, since our data is insensitive to the lower atmosphere.  A final confirmation of the abundance of CO and its vertical distribution requires further near- and mid-IR measurements of CO in the troposphere.  \\vskip 25pt  \\center{\\bf Acknowledgements}  This work was supported in part by NASA grant NAG5-7946."
    },
    "0002/astro-ph0002107_arXiv.txt": {
        "abstract": "The \\emph{ISO-SWS}\\footnote{% Based on observations made with ISO, an ESA project with instruments funded by ESA member states (especially the PI countries: France, Germany, The Netherlands and the United Kingdom) and with the participation of ISAS and NASA. The SWS is a joint project of SRON and MPE. } \\( 2.5 \\)--\\( 45\\, \\mu \\mathrm{m} \\) infrared spectroscopic observations of the nucleus of the Seyfert 2 galaxy \\( \\NGC  \\) (see companion paper) are combined with a compilation of UV to IR narrow emission line data to determine the spectral energy distribution (SED) of the obscured extreme-UV continuum that photoionizes the narrow line emitting gas in the active galactic nucleus. We search a large grid of gas cloud models and SEDs for the combination that best reproduces the observed line fluxes and NLR geometry. Our best fit model reproduces the observed line fluxes to better than a factor of 2 on average and is in general agreement with the observed NLR geometry. It has two gas components that are consistent with a clumpy distribution of dense outflowing gas in the center and a more extended distribution of less dense and more clumpy gas farther out that has no net outflow. The best fit SED has a deep trough at \\( \\sim \\! 4 \\) Ryd, which is consistent with an intrinsic Big Blue Bump that is partially absorbed by \\( \\sim \\! 6\\times 10^{19}\\, \\mathrm{cm}^{-2} \\) of neutral hydrogen interior to the NLR. ",
        "introduction": "The intrinsic spectral energy distribution (SED) of active galactic nuclei (AGN), which extends from the radio up to \\( \\gamma  \\)-rays, cannot be directly observed from the Lyman limit and up to several hundred eV due to Galactic and intrinsic absorption. However, the extreme-UV (EUV) and soft X-ray continuum can be investigated indirectly by the infrared coronal line emission. These lines are emitted by collisionally excited forbidden fine-structure transitions of highly ionized atoms, whose ionization potentials extend well beyond the Lyman limit up to hundreds of eV. Unlike the strong permitted lines of these ions, which are also emitted in the obscured EUV, the reddening-insensitive forbidden IR coronal lines and semi-forbidden optical coronal lines can be observed. Therefore, when photoionization is the main ionization mechanism, the coronal lines can provide information on the intrinsic obscured SED and the accretion process that powers the AGN. This information can be extracted by photoionization models of the NLR.  The coronal lines are collisionally suppressed in the dense broad line region (BLR) close to the continuum source and are efficiently emitted only from the more rarefied gas in the narrow line region (NLR), hundreds of pc away from the center. It is well established that large quantities of gas attenuate the continuum emission in many AGN. These gas clouds are detected by narrow UV absorption lines (e.g. Kriss et al. \\cite{Kriss92}) or by X-ray absorption features and emission lines (e.g. George et al. \\cite{George98}). Although their exact location along the line of sight is unknown, there are reasons to believe that in some cases they may be inside the NLR. In particular, the warm absorbers that block the X-ray continuum appear to cover a large fraction of the continuum source (George et al. \\cite{George98}). This raises the possibility that in some AGN the ionizing SED, which is traced by the coronal lines, is not the intrinsic one produced by the accretion process, but rather one that is filtered by intervening absorbers inside the NLR. It has been proposed that such absorbers are common in Seyfert galaxies, and are responsible for the observed correlations between the soft X-ray slope and the narrow emission line spectra of Seyfert 1.5 galaxies (Kraemer, Ruiz \\& Crenshaw \\cite{KTCG99}).  This study of the Seyfert 2 galaxy \\( \\NGC  \\) is part of the \\emph{ISO-SWS} program on bright galactic nuclei. Previous studies in this program include the reconstruction of the SED of the Seyfert 2 Circinus galaxy (Moorwood et al. \\cite{Moorwood96}; Alexander et al. \\cite{Alexander99}) and of the Seyfert 1 Galaxy NGC~4151 (Alexander et al. \\cite{Alexander99}). In both cases we found evidence of a ``Big Blue Bump'' signature of a thin accretion disk (Shakura \\& Sunyaev \\cite{SS73}). However, in the case of NGC~4151 this structure is masked by a deep absorption trough of an absorber situated between the BLR and the NLR, which filters the light that photoionizes the NLR.  In this paper we apply our SED reconstruction method to \\( \\NGC  \\), one of the closest, brightest and most extensively studied Seyfert 2 galaxies, which is considered a prototype of this AGN class. The first detection of broad permitted emission lines in the polarized light of \\( \\NGC  \\) (Antonucci \\& Miller \\cite{AM85}) provided a major argument for the Seyfert 1 and 2 unification scheme (Antonucci \\cite{Antonucci93}). This scheme postulates that the two Seyfert types have both broad and narrow line regions and an obscuring torus that lies between the two. When the torus is face on and the BLR is directly observed, the AGN is classified as a Seyfert 1. When the BLR is obscured by the torus, the AGN is classified as a Seyfert 2, and the BLR can be observed only indirectly in scattered polarized light. The factors that determine the accretion properties and the ionizing SED of AGN are currently unknown. However, the Seyfert unification picture implies that all possibly relevant factors being equal, such as luminosity, host galaxy type or redshift, the intrinsic SED of both AGN types should be similar. It is therefore of interest to complement our previous study of the nearby luminous Seyfert 1 galaxy NGC~4151 with a corresponding study of a nearby luminous Seyfert 2 galaxy with a similar host galaxy type, such as \\( \\NGC  \\). The \\emph{ISO-SWS} observations of \\emph{\\( \\NGC  \\)} are presented in a companion paper (Lutz et al. \\cite{Lutz00}) and are used there to derive the gas density and to place constraints on the structure and dynamics of the NLR.  This paper is organized as follows. In \\S\\ref{sec:physprop} we summarize the physical properties of the nucleus of \\( \\NGC  \\) that are needed for constructing the photoionization models and constraining their results. In \\S\\ref{s:lines} we present the emission line flux compilation that we use in our modeling. In \\S\\ref{s:models} we briefly discuss the construction and fitting of the NLR photoionization models. We present the results in \\S\\ref{s:results} and discuss them in \\S\\ref{s:discuss}. ",
        "conclusions": "\\label{s:discuss}  As was discussed in detail by Alexander et al. (\\cite{Alexander99}), there are various degeneracies between the parameters that describe the gas model (\\( n, \\) \\( A \\), \\( F \\), \\( U \\), \\( C \\), \\( \\theta  \\), and \\( \\tion  \\)). These degeneracies allow the fit procedure to converge to a robust best-fit SED even when the assumed values of \\( n \\), \\( A \\), \\( F \\) or \\( U \\) significantly differ from their true values, since this can be compensated to a large degree by a suitable modification of the gas geometry (\\( C \\), \\( \\theta  \\), and \\( \\tion  \\)). This property of the fit procedure is especially important in the analysis of \\( \\NGC  \\), where observations indicate that the actual properties of the NLR gas are much more complex than can be modeled by our family of simplified gas models. For this reason, we place more weight on the fact that all the best-fit SED models, whether one or two-component, display a deep trough at 4 Ryd than on the determination of the exact values of the gas parameters.  The trough in the SED is required for reproducing the relative line fluxes of the high and low ionization species. To check the robustness of this result, we attempted to re-fit the observed line fluxes with all of our one and two-component models using an approximate single power-law SED (\\( \\log f=-25.8,\\, -26.6,\\, -27.4,\\, -28.2 \\) at 2, 4, 8 and 16 Ryd, respectively), which was held fixed in the fit procedure. In all cases the low excitation lines were over-estimated with respect to the high excitation lines, regardless of the values of \\( U \\), \\( n \\) or \\( F \\). For example, when the power-law SED is applied to the best-fit two-component gas model (Table~\\ref{tbl:fit}), the low excitation lines (\\( \\Eion \\lesssim 50\\, \\mathrm{eV} \\)) are over-estimated by the model by up to a factor of 13, whereas the high excitation lines (\\( \\Eion \\gtrsim 100\\, \\mathrm{eV} \\)) are under-estimated by up to a factor of 23. The overall mismatch of this SED with the observed line fluxes is reflected in both the poor fit score of \\( S=4.1 \\) and in the very strong residual anti-correlation between the line ratios and \\( \\Eion  \\), whose random probability is \\( 10^{-4} \\). We caution against generalizing this result to mean that every AGN that exhibits a hard emission line spectrum has an absorbed ionizing SED. The emission line spectrum reflects the gas parameters, such as \\( U \\), \\( n \\) and \\( F \\), no less than it does the ionizing SED. It is necessary to have some knowledge of the likely range of values for these parameters in order to interpret the hardness of the line spectrum. Alexander et al. (\\cite{Alexander99}) provide a counter-example where subtle cancellations between the SED and the gas properties lead to a situation where an AGN (NGC4151) with a hard absorbed SED has a \\emph{softer} emission line spectrum than another AGN (Circinus) with a soft unabsorbed Big Blue Bump.  Although we do not claim to fix the gas parameters with certainty, the best fit model (Table~\\ref{tbl:fit}) is reassuringly consistent with the observations, which broadly indicate that the integrated NLR emission originates in two components. Component A of the model can be interpreted as a system of dense (\\( n=10^{4} \\)~\\( \\mathrm{cm}^{-3} \\) ), centrally concentrated (\\( 0.3\\arcsec <\\theta <1.4\\arcsec  \\)) outflowing gas clouds (\\( F\\propto r^{-2} \\)) with a relatively high filling factor (\\( F\\sim 0.01 \\)) and high ionization parameter \\( (\\log U=-1 \\)). Component B of the model can be interpreted as a more extended distribution (\\( 1.9\\arcsec <\\theta <4.5\\arcsec  \\)) of lower density gas (\\( n=2\\times 10^{3} \\)~\\( \\mathrm{cm}^{-3} \\)) with no net outflow (\\( F=\\mathrm{const}. \\)), with a lower filling factor (\\( F=0.001) \\) and a lower ionization parameter (\\( \\log U=-2 \\)). Component A contributes 58\\% of the total line flux in 22 lines listed in Table~\\ref{tbl:flux}, with the contribution to individual lines ranging from 45\\% of the low excitation \\( \\bSIVb  \\) (\\( \\Eion =34.8\\, \\mathrm{eV} \\)) to more than 99.9\\% of the very high excitation line \\( \\bSiIXbB  \\) (\\( \\Eion =303.2\\, \\mathrm{eV} \\)). Its large covering factor indicates that there is probably a significant contribution of flux from the inner \\( \\sim \\! 1\\arcsec  \\) of the diffuse SW emission cone as well as from the bright NE cone. Component B contributes the remaining 42\\% of the total line flux, mainly in the lower excitation lines. Its covering factor is small enough for it to be concentrated mostly in the NE bright emission cone, as is observed.  The best-fit procedure indicates that models with the low oxygen abundance set fit the observed line fluxes somewhat better than models with the high nitrogen abundance set. In particular, the \\( \\mathrm{O}\\, {\\textsc {iii]}}\\, \\lambda 1663 \\), whose unusual relative weakness was an important argument for assuming non-solar abundances (Netzer \\cite{Netzer97}; Kraemer et al. \\cite{KRC98}), is well reproduced by the best-fit low oxygen model with a model-to-data line flux ratio of 1.4 (Fig.~\\ref{fig:ratios}), even though it was not included in our fit since it didn't pass the minimal flux criterion.  The trough in the best fit SED (Fig.~\\ref{fig:mix}) can be interpreted as an absorption trough due to an absorber between the continuum source and the NLR. The energy resolution of the SED template, which is limited by the computational cost of enumerating over all the SED combinations, is too low to allow detailed modeling of the absorber. Figure~\\ref{fig:abs_sed} shows an example of how a quasi-thermal big blue bump that is absorbed by neutral hydrogen would appear in our low resolution SED reconstruction. We find that the trough is consistent, for example, with an absorber that shadows the entire NLR (\\( C_{\\mathrm{abs}}=1) \\) and has a column density of \\( N_{H^{0}}=6\\times 10^{19}\\, \\mathrm{cm}^{-2} \\) in neutral hydrogen, or with an absorber that allows a small leakage of unfiltered radiation (\\( C_{\\mathrm{abs}}=0.999 \\)) and a column density of \\( N_{H^{0}}=10^{20}\\, \\mathrm{cm}^{-2} \\). A similar trough, consistent with an absorber of \\( C_{\\mathrm{abs}}>0.99 \\) and \\( N_{H^{0}}=5\\times 10^{19}\\, \\mathrm{cm}^{-2} \\), was discovered in the reconstructed SED of Seyfert 1 galaxy NGC~4151 (Alexander et al. \\cite{Alexander99}). That absorber was also detected in the \\emph{HUT} absorption line spectra of the UV continuum of NGC~4151 (Kriss et al. \\cite{Kriss92}, \\cite{Kriss95}) . The \\emph{HUT} spectra of \\( \\NGC  \\) (Kriss et al \\cite{KDBFL92}) do not have a high enough S/N to allow the detection of absorption lines in the scattered UV continuum of this AGN or against the stellar background (G. Kriss, private comm.). We predict that future sensitive absorption line studies should reveal the presence of such an absorber.  The bias in our results due to the fact that we neglected the line emission from the absorbing gas is likely to be small if the absorber is similar to the dense, high velocity UV absorber that was detected in NGC~4151. Such an absorber will not emit forbidden lines, and its permitted lines will be broader than typical NLR lines. Only 3 of the 22 lines we used in our fit are permitted lines, and we did not use lines that are contaminated by broad components. A highly ionized and optically thin absorber will produce strong \\( \\OVI  \\) line emission in excess of the typical NLR emission. It is therefore interesting that unlike the two forbidden \\( \\bOIIIbB  \\) and \\( \\bOIVb  \\) lines and the semi-forbidden \\( \\OIIIb  \\) line, which are well reproduced by the best fit model, the observed \\( \\OVI  \\) line is 3.2 stronger than predicted (Fig.~\\ref{fig:ratios}).  We have, up to now, applied our SED reconstruction method to \\emph{ISO-SWS} observations of IR coronal lines of three AGN: the Seyfert 2 Circinus galaxy (Moorwood et al. \\cite{Moorwood96}; Alexander et al. \\cite{Alexander99}), the Seyfert 1 galaxy NGC~4151 (Alexander et al. \\cite{Alexander99}), and the Seyfert 2 galaxy \\( \\NGC  \\) (this work). In one of these (Circinus), we detect a Big Blue Bump that peaks at \\( \\gtrsim 50\\, \\mathrm{eV} \\). In the other two we detect deep troughs, which are consistent (but not exclusively so) with a Big Blue Bump that is absorbed by neutral gas interior to the NLR. Our findings thus far are consistent with the picture that luminous Seyfert galaxies are powered by thin accretion disks that produce a quasi-thermal Big Blue Bump, and that in a large fraction of them the NLR sees a partially absorbed ionizing continuum, as suggested by Kraemer et al. (\\cite{KRC98})."
    },
    "0002/astro-ph0002113_arXiv.txt": {
        "abstract": "A growing number of astronomical resources and data or information services are made available through the Internet. However valuable information is frequently hidden in a deluge of non-pertinent or non up-to-date documents. At a first level, compilations of astronomical resources provide help for selecting relevant sites. Combining yellow-page services and meta-databases of active pointers may be an efficient solution to the data retrieval problem. Responses generated by submission of queries to a set of heterogeneous resources are difficult to merge or cross-match, because different data providers generally use different data formats: new endeavors are under way to tackle this problem. We review the technical challenges involved in trying to provide general search and discovery tools, and to integrate them through upper level interfaces. ",
        "introduction": "How to help the users find their way through the jungle of information services is a question which has been raised since the early development of the WWW (see e.g., Egret \\cite{jungle}), when it became clear that a big  centralized system was not the efficient way to go.  Obviously the World Wide Web is a very powerful medium for the development of distributed resources: on the one hand the WWW provides a common medium for all information providers -- the language is flexible enough so that it does not bring unbearable constraints on existing databases -- on the other hand the distributed hypertextual approach opens the way to navigation and links between services (provided a  minimum of coordination can be achieved). Let us note that it has been already widely demonstrated that coordinating spirit is not out of reach in a small community such as astronomy, largely sheltered from commercial influence.  Searching for a resource (either already visited, or unknown but  expected), or browsing lists of existing services in order to discover new tools of interest implies a need for query strategies that cannot generally be managed at the level of a single data provider.  There is a need for road-guides pointing to the most useful resources, or to compilations or databases where information can be found about these resources. Such guides have been made in the past, and are of very practical help for the novice as well as the  trained user, for example: Andernach et al.\\ \\cite{AHM94}, Egret \\& Heck \\cite{waw}, Egret \\& Albrecht \\cite{amp2}, Heck \\cite{eppa}, Grothkopf \\cite{librarians}, Andernach \\cite{ALDIA}.  In the present paper our aim is to address the questions related to the collection, integration and interfacing of the wealth of astronomical Internet resources, and also to describe some strategies that have to be developed for building cooperative tools which will be essential in the research environment of the decade to come. ",
        "conclusions": "The on-line ``Virtual Observatory'' is currently under construction with on-line archives and services potentially giving access to a huge quantity of scientific information: its services will allow astronomers to select the information of interest for their research, and to access original data, observatory archives and results published in journals. Search and discovery tools currently in development will be of vital importance to make all the observational data and information available to the widest scientific community."
    },
    "0002/astro-ph0002439_arXiv.txt": {
        "abstract": "The Local Group irregular galaxy IC 1613 has remained an enigma for many years because of its apparent lack of star clusters. We report the successful search for clusters among several of the candidate objects identified many years ago on photographic plates. We have used a single HST WFPC2 pointing and a series of  images obtained with the WIYN telescope under exceptional seeing conditions, examining a total of 23 of the previously published candidates. All but six of these objects were found to be either asterisms or background galaxies. Five of the six remaining candidates possibly are  small, sparse clusters and the sixth, C32, is an obvious cluster. It is a compact, young object, with an age of less than 10 million years and a total absolute magnitude of M$_V = -5.78\\pm0.16$ within a radius of 13 pc. ",
        "introduction": "\\citet{ba63} remarked on the fact, which he obviously considered remarkable, that the irregular galaxy IC 1613, a member of the Local Group that he had studied quite intensively, appeared to be devoid of star clusters. This fact was again discussed by \\citet{va79}, who compared IC 1613 with the SMC, pointing out that the SMC's many rich clusters are so conspicuous that, if IC 1613 has any, they should show up clearly on available plate material. \\citet{ho78} had searched photographic plates taken with the Palomar Observatory's 5m and the Lick Observatory's 3m telescopes and published a list of 43 possible candidates for clusters, all of which were very inconspicuous, none having more than 6 resolved stars on the best images. Clearly, these objects were not comparable to the rich star clusters of the Magellanic Clouds, but he thought that at least some of them might be similar to small Galactic open clusters like the Pleiades. More recently \\citet{fr88}, in her CCD-based study of the color-magnitude diagram (CMD) of a portion of IC 1613, noted that she could not identify those candidate clusters that should have been visible on her images. On the other hand, \\citet{ge99} observed  a field centered on the most active area of star formation (east of our present field), identified 12 of the earlier cluster candidates and measured integrated UBV magnitudes for eight, six of which they found to have colors indicating ages of less than 10 million years. They also identified two new small, young clusters. They searched for massive clusters, either young or old globular clusters, but found none. For a related discsussion of the young massive cluster populations in galaxies and their correlations with integrated galaxy properties see \\citet{la00}.  Thus it has been  clear that IC 1613 must be poor in rich star clusters compared with the Magellanic Clouds, a fact which van den Bergh attributed to the lack of a history of strong shocks. But it has not been clear whether or not IC 1613 is lacking in a significant number of normal open clusters, as most of the published candidate clusters have not been examined on anything but the original photographic plates. This paper reports the results of a study of some of these candidates using HST and WIYN\\footnote{The WIYN Observatory is a joint facility of the University of Wisconsin-Madison, Indiana University, Yale University, and the National Optical Astronomy Observatories.} telescope images. ",
        "conclusions": "We have surveyed the central area of IC 1613 for star clusters. There are no globular clusters or massive young clusters present, but there are six very sparse groupings that may be open clusters like some of the smaller ones in our Galaxy. One of them, previously cataloged as C32, is a very young cluster embedded in an HII region and surrounded by an OB association.  When the uncertain nature of many of the cluster candidates, their very small sizes and their implied small stellar populations are taken into account, it is clear that IC 1613 is a cluster-poor galaxy."
    },
    "0002/astro-ph0002325_arXiv.txt": {
        "abstract": "We analyse a sample of 236 \\C\\ from the \\HP\\ catalog, using the method of ``reduced parallaxes'' in $V, I, K$ and the reddening-free ``Wesenheit-index''. We compare our sample to those considered by Feast \\& Catchpole (1997) and Lanoix et al. (1999), and argue that our sample is the most carefully selected one with respect to completeness, the flagging of overtone pulsators, and the removal of \\C\\ that may influence the analyses for various reasons (double-mode \\C, unreliable \\HP\\ solutions, possible contaminated photometry due to binary companions).  From numerical simulations, and confirmed by the observed parallax distribution, we derive a (vertical) scale height of \\C\\ of 70 pc, as expected for a population of 3-10 \\msol\\ stars. This has consequences for Malmquist- and Lutz-Kelker (Lutz \\& Kelker 1973, Oudmaijer et al. 1998) type corrections which are smaller for a disk population than for a spherical population.   The $V$ and $I$ data suggest that the slope of the Galactic $PL$-relations may be shallower than that observed for LMC Cepheids, either for the whole period range, or that there is a break at short periods (near $\\log P_0 \\approx 0.7-0.8$).  We stress the importance of two systematic effects which influence the distance to the LMC: the slopes of the Galactic $PL$-relations and metallicity corrections. In order to assess the influence of these various effects, we present 27 distance moduli (DM) to the LMC. These are based on three different colours ($V,I,K$), three different slopes (the slope observed for \\C\\ in the LMC, a shallower slope predicted from one set of theoretical models, and a steeper slope as derived for Galactic \\C\\ from the surface-brightness technique), and three different metallicity corrections (no correction as predicted by one set of theoretical models, one implying larger DM as predicted by another set of theoretical models, and one implying shorter DM based on empirical evidence). We derive DM between 18.45 $\\pm$ 0.18 and 18.86 $\\pm$ 0.12. The DM based on $K$ are shorter than those based on $V$ and $I$ and range from 18.45 $\\pm$ 0.18 to 18.62 $\\pm$ 0.19, but the DM in $K$ could be systematically too low by about 0.1 magnitude because of a bias due to the fact that NIR photometry is available only for a limited number of stars.  From the Wesenheit-index we derive a DM of 18.60 $\\pm$ 0.11, assuming the observed slope of LMC \\C\\ and no metallicity correction, for want of more information.   The DM to the LMC based on the parallax data can be summarised as follows. Based on the $PL$-relation in $V$ and $I$, and the Wesenheit-index, the DM is \\begin{displaymath} 18.60 \\pm 0.11 \\;\\; (\\pm 0.08 \\;{\\rm slope})(^{+0.08}_{-0.15} \\;{\\rm metallicity}), \\end{displaymath} which is our current best estimate. Based on the $PL$-relation in $K$ the DM is $\\;\\;\\;\\; 18.52 \\pm 0.18$ \\begin{displaymath} \\;\\;(\\pm 0.03 \\;{\\rm slope}) (\\pm 0.06 \\;{\\rm metallicity}) (^{+0.10}_{-0} \\;{\\rm sampling \\;bias}). \\end{displaymath} The random error is mostly due to the given accuracy of the \\HP\\ parallaxes and the number of Cepheids in the respective samples.  The terms between parentheses indicate the possible systematic uncertainties due to the slope of the Galactic $PL$-relations, the metallicity corrections, and in the $K$-band, due to the limited number of stars. Recent work by Sandage et al. (1999) indicates that the effect of metallicity towards shorter distances may be smaller in $V$ and $I$ than indicated here.  From this, we point out the importance of obtaining NIR photometry for more (closeby) \\C, as for the moment NIR photometry is only available for 27\\% of the total sample. This would eliminate the possible bias due to the limited number of stars, and would reduce the random error estimate from 0.18 to about 0.10 mag. Furthermore, the sensitivity of the DM to reddening, metallicity correction and slope are smallest in the $K$-band. ",
        "introduction": "Cepheids are important standard candles in determining the extra-galactic distance scale. The results of the \\HP\\ mission allow, in principle, a calibration of the period-luminosity relation based on the available parallaxes. Feast \\& Catchpole (1997; hereafter FC) did just that for the $M_{\\rm V} - \\log P$-relation based on pre-released \\HP\\ data of 223 Cepheids available to them at that time. Now that the entire catalog has become available (ESA 1997) it is timely to analyse the full sample of Cepheids in it.  In a recent paper, Lanoix et al. (1999, hereafter L99) presented a study similar to ours and they derived the zero points of the $M_{\\rm V} - \\log P$- and $M_{\\rm I} - \\log P$-relations, without, however, discussing the distance to the LMC. We will indicate where the two studies agree and differ.  The paper is organised as follows. In Sect.~2 the sample considered in the present paper is presented, and compared to that in FC and L99. In Sect.~3 several aspects involved in the analysis of parallax data are described, and the method of ``reduced parallaxes'' is outlined, together with all necessary recipes to obtain the reddening. In Sect.~4 the zero points of the $PL$-relations in $V,I, K$ and the reddening-free ``Wesenheit-index'' (e.g. Tanvir 1999, and Eq.~(11)) are presented for different selections of the sample, which are discussed in Sect.~5. In Sect.~6 we construct and present the zero points for volume complete samples of stars. In Sect.~7 we describe numerical simulations that are first of all tuned to fit the observed properties of the \\C\\ in the \\HP\\ catalog, and then are used to show that the method of ``reduced parallaxes'' introduces a bias which is of the order of 0.01 mag or less.  Based on these results we discuss in Sect.~8 the distance to the LMC, and elaborate on the various uncertainties.  \\\\ ",
        "conclusions": "\\subsection{On the use of longer wavelengths}  One argument to consider $PL$-relations in $I$ and $K$ besides the traditional relation in $V$, is because of the smaller extinction at longer wavelengths. We have calculated the zero points for the whole sample considering a systematic shift for all stars of 0.005 mag in $V$ (respectively $I$ and $K$), 0.007 in $B-V$ (resp. $V-I$, $J-K$), a shift in the period-colour relations (Eqs.~(4), (8), (12)) equal to 1/10-th of the quoted dispersions, and a shift in the selective reddening (Eqs.~(5), (10), (14)) equal to the quoted dispersion. The results are listed in Table~1 (solutions 48-51), Table~2 (solutions 23-26), and Table~3 (solutions 15-18). Adding the differences between these zero points and that for the default case in quadrature, one arrives at estimated uncertainties due to errors in the photometry, reddening and period-colour relations of 0.038 mag in $V$, 0.032 mag in $I$, and 0.011 mag in $K$. This illustrates that the $K$-band $PL$-relation is indeed the least sensitive to these effects. For the Wesenheit-index the error due to the photometry and reddening coefficients amounts to 0.028 mag (solutions 13-15 in Table~4). \\\\  Unfortunately, the advantages of using the infrared, like the intrinsically tighter $PL$-relation and the insensitivity to reddening, are countered by the fact that so few stars have been measured in the NIR so far. Of the approximately 48 stars with in 1 kpc, $BV$ photometry is available for all of them, $VI$ for 41 of them, but $JHK$ data (of sufficient quality) for only 24. It is estimated that determining the intensity-mean NIR magnitudes of the remaining 24 stars alone would bring the scatter in the zero point determination in the $K$-band down from about 0.17 to less than 0.1 mag. \\\\   \\subsection{On the slopes of Galactic $PL$-relations}  Another important issue concerns the slopes in the respective Galactic $PL$-relations. Common practice is to adopt the slope determined for \\C\\ in the LMC, but the slope could be different for Galactic \\C.  In Table~5 we have collected slopes of the $PL$-relations from the literature for \\C\\ in the Galaxy, LMC and SMC, in $V,I,K$, both observationally determined and from two recent theoretical papers (Alibert et al. 1999, Bono et al. 1999 and private communication). Table~5 includes ths slopes in $V$ and $I$ by Madore \\& Freedman (1991) used by the {\\sc hst} $H_0$ Key Project (see for example Gibson et al. 1999).  From the results of Gieren et al. (1998) on the \\C\\ in the LMC, we have calculated the error in the slope using the data they kindly provided (their Tables 8 and 9). In addition we have calculated the slope and zero point for their data set but using cut-offs in period, for reasons that will be explained later.  \\\\  There are several interesting features to be noted about the slopes. Observationally, the slopes in the $V$ and $I$ $PL$-relations in the LMC are very well established, respectively, about $-2.81$ and $-3.05$, and these are the default slopes adopted in the present study, with errors of about 0.08 and 0.06. For $M_{\\rm K}-\\log P$-relation there are fewer observational data available but the slope in Gieren et al. (1998) is determined very accurately and is in reasonable agreement with the result of Madore \\& Freedman (1991). \\\\   One fact needs to be mentioned however, and that is that the period distribution of the calibrating \\C\\ in the LMC is very different from that of the Galactic \\C\\ in the \\HP\\ catalog. In Table 8 in Gieren et al. (1998) there are 53 LMC \\C\\ listed with $V,I$ photometry that define their $PL$-relation.  Nine have $\\log P \\le 0.555$, then there is a gap, and 42 have $\\log P \\ge 0.846$. In $JHK$ (their Table 9), there are 59 \\C, including 5 that have $\\log P \\le 0.680$, then there is a gap, and 54 have $\\log P \\ge 0.834$.  The same dichotomy of \\C\\ in period can be seen from Tanvir (1999). Note that in more distant Galaxies than the Magellanic Clouds, due to observational bias, most known Cepheids have periods longer than 10 days. We therefore have included in Table~5 the slopes in $V,I,K$ based on the data in Gieren et al. (1998), but have divided the sample according to period. The slopes for the long period sub-sample differ only slightly from that for the whole sample but are characterised by slightly larger errors. The slopes for the short-period sample have larger errors because of the small number of stars involved but nevertheless are systematically steeper at the 1$\\sigma$ level in all three colours. For comparison, in our sample only 107 of the 236 \\C\\ have $\\log P \\ge 0.846$, and the zero point of this sample differs at the 1.4$\\sigma$ level from that of the whole sample (solution 52 in Table~1). This is yet another indication that for the default slope of $-2.81$ the zero point may depend on the period range chosen. For a slope of $\\delta = -2.22$ (compare solutions 38, 53) this is not the case. \\\\  A second issue concerns the predictions of the theoretical models and the comparison with observations. In $V,K$ for the LMC, and in $I$ for the SMC the models of Bono et al. (1999) are in very good agreement with the observed slope, in $I$ for the LMC and $V$ for the SMC the agreement is poor. In $I, K$ for the LMC, the models of Alibert et al. (1999) are in good agreement with the observed slopes; in $V$ the agreement for the LMC and $V,I$ for the SMC is fair. However, the prediction both models make for the Galaxy are very different. Gieren et al. (1998) have derived $PL$-relations for Galactic \\C\\ using the surface brightness technique. The slopes they find in all three bands are {\\it steeper} than the corresponding slopes in the LMC at the 2-3 $\\sigma$ level. In their paper they ascribe this to small number statistics, and in the end assume the LMC slopes to hold for the Galactic \\C\\ as well, also, they add, because the slopes for the LMC \\C\\ are better determined.  The models of Alibert et al. (1999) predict slopes for the Galactic $PL$-relations that are {\\it steeper} than the observed ones in the LMC in $V,I,K$ as well, although by a small amount only (the main conclusion of their paper is actually that the slope and zero point of the $PL$-relations do not depend on metallicity). By contrast, the Bono et al. (1999) paper predicts slopes that are significantly {\\it shallower} in the Galaxy compared to the LMC especially in $V,I$ but also in $K$. In addition, the Bono et al. models actually predict a non-linear $PL$-relation in $V$, for all metallicities considered.  Furthermore, Alibert et al. mention that a change of slope in $PL$-relations is expected at short periods due to the reduction of the blue loop during core He burning, and that this change of slope occurs near $\\log P = 0.2$ for $Z$ = 0.004, $\\log P = 0.35$ for $Z$ = 0.01, and $\\log P = 0.5$ for $Z$ = 0.02. Such a change of slope (in the sense that the slope of the $PL$-relation is steeper for periods below this limit) was recently observed for SMC \\C\\ (Bauer et al. 1999) with the break occurring near $\\log P = 0.3$, near the predicted value. \\\\    \\begin{table*} \\caption[]{Slopes of the $PL$-relations.} \\begin{tabular}{ccll} \\hline Slope         & Colour & System & Reference \\\\ \\hline $-3.037 \\pm 0.138$ & V & GAL & Gieren et al. (1998)\\\\ $-2.22 \\pm 0.04$  & V & 0.02& Bono et al. (1999); non-linear, $\\log P<1.4$\\\\ $-2.905$           & V & 0.02& Alibert et al. (1999)  \\\\  $-2.810 \\pm 0.082$ & V & LMC & Tanvir (1999) \\\\ $-2.769 \\pm 0.073$ & V & LMC & Gieren et al. (1998) \\\\ $-2.820 \\pm 0.118$ & V & LMC & this work; 44 stars with $\\log P > 0.845$ from Gieren et al. (1998)\\\\ $-3.54 \\pm 0.68$  & V & LMC & this work;  9 stars with $\\log P < 0.845$ from Gieren et al. (1998)\\\\ $-2.88 \\pm 0.20$  & V & LMC & Madore \\& Freedman (1991)\\\\ $-2.81 \\pm 0.06$  & V & LMC & Caldwell \\& Laney (1991)\\\\ $-2.79 \\pm 0.06$  & V & 0.008& Bono et al. (1999); non-linear, $\\log P<1.4$\\\\ $-2.951$           & V & 0.01& Alibert et al. (1999)  \\\\  $-2.63 \\pm 0.08$   & V & SMC & Caldwell \\& Laney (1991), \\\\ $-3.04 \\pm 0.04$   & V & 0.004& Bono et al. (1999); non-linear, $\\log P<1.4$\\\\ $-2.939$           & V & 0.004& Alibert et al. (1999)  \\\\  $-3.329 \\pm 0.132$ & I & GAL & Gieren et al. (1998)\\\\ $-2.35 \\pm 0.08$  & I & 0.02& Bono et al. (1999)  \\\\ $-3.102$           & I & 0.02& Alibert et al. (1999)\\\\  $-3.078 \\pm 0.059$ & I & LMC & Tanvir (1999) \\\\ $-3.041 \\pm 0.054$ & I & LMC & Gieren et al. (1998)\\\\ $-3.084 \\pm 0.088$ & I & LMC & this work; 44 stars with $\\log P > 0.845$ from Gieren et al. (1998)\\\\ $-3.39 \\pm 0.39$  & I & LMC & this work;  9 stars with $\\log P < 0.845$ from Gieren et al. (1998)\\\\ $-3.14 \\pm 0.17$  & I & LMC & Madore \\& Freedman (1991)\\\\ $-3.01 \\pm 0.05$  & I & LMC & Caldwell \\& Laney (1991) \\\\ $-2.63 \\pm 0.08$  & I & 0.008& Bono et al. (1999)  \\\\ $-3.140$          & I & 0.01& Alibert et al. (1999) \\\\  $-2.92 \\pm 0.07$  & I & SMC & Caldwell \\& Laney (1991)  \\\\ $-2.73 \\pm 0.10$  & I & 0.004& Bono et al. (1999)  \\\\ $-3.124$          & I & 0.004& Alibert et al. (1999) \\\\  $-3.598 \\pm 0.114$ & K & GAL & Gieren et al. (1998) \\\\ $-3.03 \\pm 0.07$  & K & 0.02 & Bono et al. (1999)  \\\\ $-3.367$          & K & 0.02 & Alibert et al. (1999)\\\\  $-3.267 \\pm 0.041$ & K & LMC & Gieren et al. (1998) \\\\ $-3.304 \\pm 0.052$ & K & LMC & this work; 54 stars with $\\log P > 0.833$ from Gieren et al. (1998)\\\\ $-3.37 \\pm 0.39$  & K & LMC & this work;  5 stars with $\\log P < 0.833$ from Gieren et al. (1998)\\\\ $-3.42 \\pm 0.09$  & K & LMC & Madore \\& Freedman (1991) \\\\ $-3.19 \\pm 0.09$  & K & 0.008 & Bono et al. (1999) \\\\ $-3.395$          & K & 0.01& Alibert et al. (1999) \\\\  $-3.27 \\pm 0.09$  & K & 0.004 & Bono et al. (1999)  \\\\ $-3.369$          & K & 0.004 & Alibert et al. (1999) \\\\  $-3.411 \\pm 0.036$ & Wesenheit & LMC & Tanvir (1999) \\\\   \\hline \\end{tabular} \\end{table*}    \\subsection{Consistency between individual distances based on different colours}  For a given slope and zero point of the $PL$-relation one can calculate the photometric distance, and since we have determined the zero point for three photometric band we can intercompare photometric distances to individual \\C. This is illustrated in Fig.~6, for the default slopes of the $PL$-relations. The zero point of the $M_{\\rm V}-\\log P$ relation is fixed at $-1.411$ (solution 10, Table~1). The zero point of the $M_{\\rm I}-\\log P$ relation is determined to give a mean difference in $(m-M)_{\\rm 0,I} - (m-M)_{\\rm 0,V}$ of zero, and is found to be $-1.918$ (top panel Fig.~6). The rms dispersion is 0.14 mag.  Similarly, to create the bottom panel, the zero point of the $M_{\\rm K}-\\log P$ relation was determined to give a mean difference in $(m-M)_{\\rm 0,K} - (m-M)_{\\rm 0,I}$ of zero, and is found to be $-2.600$, with an rms dispersion is 0.10. The values for the zero points in $I$ and $K$ derived in this way differ only slightly from the solutions 1 in Tables 2 and 3.  Interestingly, least-square fitting shows that the slopes of the relations are not unity, but 0.980 $\\pm$ 0.007 (top panel), and 0.965 $\\pm$ 0.007 (bottom panel).  Using the same procedure, but adopting steeper slopes of $-3.04$, $-3.33$, $-3.60$ (see Sect.~8 for the reason of these choices) in, respectively, the $M_{\\rm V}-\\log P$, $M_{\\rm I}-\\log P$ and $M_{\\rm K}-\\log P$ relations and fixing the zero point of the $M_{\\rm V}-\\log P$ relation at $-1.234$ we find in the same way the zero points of the $M_{\\rm I}-\\log P$ and $M_{\\rm K}-\\log P$-relations to be $-1.696$ and $-2.328$. The slopes are 0.983 $\\pm$ 0.007 and 0.971 $\\pm$ 0.007.  Using the same procedure, but adopting shallower slopes of $-2.22$, $-2.35$, $-3.05$ in, respectively, the $M_{\\rm V}-\\log P$, $M_{\\rm I}-\\log P$ and $M_{\\rm K}-\\log P$ relations and fixing the zero point of the $M_{\\rm V}-\\log P$ relation at $-1.885$ (solution 38 in Table~1) we find the zero point of the $M_{\\rm I}-\\log P$ and $M_{\\rm K}-\\log P$-relations to be $-2.491$ and $-2.692$. The slopes are 0.974 $\\pm$ 0.007 and 1.010 $\\pm$ 0.010.  The conclusion is the the photometric distances based on different $PL$-relations are consistent with each other at a level of 0.10-0.14 mag, similar to the uncertainties in the individually derived zero points in $V,I,K$ for the full sample. The data does not allow to discriminate between different choices of the slopes of the $PL$-relations.   \\begin{figure} \\centerline{\\psfig{figure=vergelijk.ps,width=8.8cm}} \\caption[]{Comparison of the distance moduli based on the $PL$-relations in $V$ and $I$ (top panel), assuming zero points of respectively $-1.411$ and $-1.918$, and the default slopes, and $I$ and $K$ (bottom panel) for zero points of $-1.918$ and $-2.600$, and the default slopes. The solid line is the one-to-one relation, although a least-square fit indicates that the best fitting slope differs slightly from unity (see text). } \\end{figure}       \\subsection{The finally adopted zero points}  Based on the results obtained in Sects.~4, 5 and 6, we will present three sets of solutions, a `traditional' one, and two alternatives. The traditional one follows FC and L99 closely. The zero point adopted is the one for the entire sample (which has the lowest error of the samples that are not selected on observed parallax), adopting the slope of the $PL$-relation observed for \\C\\ in the LMC. This would then be solution 10 from Table~1 ($\\rho = -1.41 \\pm 0.10$), solution 1 from Table~2 ($\\rho = -1.89 \\pm 0.11$) and solution 1 from Table~3 corrected for the off-set in the corresponding solution in $V$ and $I$ as discussed in Sect.~4.3 ($\\rho = -2.50 \\pm 0.17$). These zero points have to be corrected for Malmquist bias. If the Malmquist bias would be evaluated in magnitude space, this bias would amount to 0.0092, 0.021 and 0.013 magnitude in $V,I,K$, respectively, for a disk population and the adopted values for ${\\sigma}_{\\rm H}$ of, respectively, 0.10, 0.15 and 0.12 mag (Stobie et al. 1989). However, if evaluated using the reduced parallax method the Malmquist bias is smaller (see Oudmaijer et al. 1999), and from the numerical simulations (Table~6) we find a Malmquist bias of 0.007 in $V$, 0.003 in $I$, and an unphysical negative value in $K$, probably due to the smaller number of stars involved. For the Malmquist bias we will assume a (round number) of 0.01 mag in all three bands.   Also, we have increased the errors in the zero points to reflect the sensitivity to uncertainties in the photometry, reddening and period-colour relations (see Sect.~5.1). For adopted slopes of $-2.81$, $-3.05$ and $-3.27$, the finally adopted zero points of the $PL$-relations in $V$, $I$ and $K$ are, respectively, $\\rho = -1.40 \\pm 0.11$, $-1.88 \\pm 0.12$ and $-2.49 \\pm 0.17$. For the Wesenheit-index, after correcting for Malmquist bias and increasing the error bar due to the uncertainties described in Sect.~5.1, we adopt a zero point of $-2.55 \\pm 0.11$ for a slope of $-3.411$. \\\\  The `traditional' solution above has certain advantages, like the fact that the errors are smallest compared to solutions that do not include all stars, or, since the slope adopted is the one in the LMC, a distance determination to the LMC is essentially a comparison of zero points only (apart from systematic effects). On the other hand, alternative solutions can be presented which also have merits. These solutions are based on a volume complete sample (at least in $V$ and $I$) to avoid Malmquist bias. Although small, its value does depend on the distribution of stars, and the intrinsic spread in the $PL$-relation and selecting a volume complete sample avoids Malmquist bias outright.  In $K$, no volume complete sample could be constructed and we used the full sample instead.  The alternative solutions take into account the theoretical prediction that the slope of the $PL$-relation may change at $\\log P \\sim 0.5$ for solar metallicities (Alibert et al. 1999).  Where the two alternative solutions differ are in the adopted slopes for the Galactic $PL$-relations. One alternative solution takes into account the, admittedly at the 1$\\sigma$ level of significance, evidence presented in Sects.~4.1-4.2 that the slope of the Galactic $PL$-relations are shallower than the ones in the LMC, and in accordance we have adopted the theoretical slopes predicted by Bono et al. (1999) for solar metallicities. The second alternative method adopts the slopes in Gieren et al. (1998), who derived distances from the infrared version of the surface brightness technique to Galactic \\C, which yields steeper slopes than for the LMC.  These alternative solutions are presented in Tables~1-3 (solutions 54-55, 28-28 and 19-20, respectively).  The zero points in $K$ have to be corrected for Malmquist bias (+0.01 mag) and for the off-sets in the corresponding $V$ and $I$ solutions, and the errors in all three zero points are increased for reasons indicated above.  For the adopted slopes of $-2.22$, $-2.35$ and $-3.05$, the zero points of the $PL$-relations in $V$, $I$ and $K$ are, respectively, $\\rho = -1.95 \\pm 0.12$, $-2.47 \\pm 0.15$ and $-2.68 \\pm 0.18$. For the adopted slopes of $-3.04$, $-3.33$ and $-3.60$, the zero points of the $PL$-relations in $V$, $I$ and $K$ are, respectively, $\\rho = -1.30 \\pm 0.13$, $-1.75 \\pm 0.14$ and $-2.30 \\pm 0.18$. \\\\   \\subsection{$PL$-relation of LMC \\C}  Before proceeding we have to adopt $PL$-relations for the LMC \\C. In $V$, for a slope of $-2.81$, Caldwell \\& Laney (1991) find a zero point of 17.23 $\\pm$ 0.02. Tanvir (1999), for the same slope, gives an observed zero point of 17.451 $\\pm$ 0.043. Adopting a mean reddening to the LMC of $E(B-V)$ of 0.08 (Caldwell \\& Laney 1991) and a ratio of total-to-selective reddening of $R_{\\rm V} = A_{\\rm V}/E(B-V) = $ 3.27 (Eq.~(5) for typical colors) we derive a de-reddened zero point of 17.19 $\\pm$ 0.04. For the \\C\\ in the LMC we adopt the weighted mean of these two values, or a zero point of 17.22 $\\pm$ 0.02 in $V$ for a slope of $-2.81$.  In $I$, for a slope of $-3.041$, Gieren et al. (1998) find a zero point of 16.74 $\\pm$ 0.06 (the error is calculated by us, from their data). Tanvir (1999), for a slope of $-3.078$, gives an observed zero point of 16.904 $\\pm$ 0.031. Adopting $A_{\\rm I} = 0.69 A_{\\rm V} = 0.18$ mag for the reddening in $V$ calculated as above, we derive a zero point of 16.72 $\\pm$ 0.03. For the default slope of $-3.05$ we adopt the weighted mean of these two values, or a de-reddened zero point in the $I$ band for the \\C\\ in the LMC of 16.72 $\\pm$ 0.02.  In $K$, for a slope of $-3.267$, Gieren et al. (1998) find a zero point of 16.03 $\\pm$ 0.05 (the error is calculated by us from their data) for the \\C\\ in the LMC. This value is adopted by us.  For the Wesenheit-index, for a slope of $-3.411$, Tanvir (1999) finds a zero point of 16.051 $\\pm$ 0.017. This value is adopted by us.\\\\  \\subsection{Metallicity correction}  We will now consider the effect of metallicity on the zero point. For comparison, FC applied a correction of +0.042 mag to the zero point in the $V$-band, based on Laney \\& Stobie (1994). The theoretical models of Bono et al. (1999), and Alibert et al. (1999) provide $PL$-relations and from those the difference $\\Delta M = M$(Gal) $- M$(LMC) can be determined which will depend on the photometric band and period. We have determined this difference for two periods, namely for $\\log P_0 = 0.77$ which we have determined to be the mean period of the volume complete sample of Galactic \\C\\ in \\HP\\, and for $\\log P_0 = 0.47$ which is the mean period of \\C\\ in the LMC (Alcock et al. 1999).  The results for $\\Delta M$ are listed in Table~7 for the three photometric bands. This illustrates the difference between the two theoretical models, for the Alibert et al. (1999) models predict essentially no dependence on metallicity, while the Bono et al. (1999) models predict a significant metallicity dependence, which mostly is in the sense that the metal-rich pulsators are fainter than the metal-poor ones. This is at variance with various empirical estimates that give the opposite result, and that, in the $V$-band, vary between $-0.24 \\pm 0.16$ (Kennicutt et al. 1998) and about $-0.4$ mag/dex (Kochanek 1997, Sasselov et al. 1997, Storm et al. 1999\\footnote{Recently, Storm (2000) suggested that this result may not be confirmed from his latest analysis and that the correction may have a positive sign instead.}). A mean of these four determinations is $-0.38 \\pm 0.09$ mag/dex, which for a difference in metallicity of 0.4 dex, implies $\\Delta M = -0.15 \\pm 0.04$ in the $V$-band. In the $I$-band we assume the same value following Sasselov et al. (1997) and Kochanek (1997), but the reader should realise that this number is less well established than the correction in $V$, and in the $K$-band adopt $\\Delta M = -0.07$.  Note however, that a metallicity dependence as large as 0.4 mag in $V$ as suggested by Sekiguchi \\& Fukugita (1998) can be excluded at the 6$\\sigma$ level (Laney 1999, 2000). In recent papers, Saio \\& Gautschy (1998) and Sandage et al. (1999) found no significant metallicity dependence on the bolometric $PL$-relation, and slopes of $-0.08$ mag/dex in $V$ and $-0.1$ mag/dex in $I$ (Sandage et al. 1999), which represents a shallower dependence than the values listed above, that are adopted in the present study, and which therefore may be an extreme view.\\\\  \\subsection{The distance to the LMC}  In Table~8 are listed the true distance moduli (DM) to the LMC for $V,I,K$, the three slopes (`traditional' meaning the observed slopes for the \\C\\ in the LMC, `shallower' adopting the theoretical slopes from Bono et al. and `steeper' adopting the observed slopes for Galactic \\C\\ from Gieren et al. (1998) and the three metallicity corrections (`0' means no correction, `+' means a longer distance scale as implied by the models from Bono et al., and `$-$' means a shorter distance scale as implied from empirical evidence). The error quoted includes the error in the zero point of both the Galactic and LMC Cepheids, and where appropriate, the error due to the metallicity correction, and for the solutions with either steeper or shallower slope, the uncertainty due to the difference in DM at $\\log P = 0.47$ and 0.77.  Also included are the weighted mean DM, averaged over $V,I,K$ (with internal error), and, for reference, the (unweighted) mean DM per photometric band of all the solutions (with the one-sided range in the solutions).  Also included is the solution based on the Wesenheit-index assuming the slope observed in the LMC, and no metallicity correction. The DM range from 18.45 $\\pm$ 0.18 to 18.86 $\\pm$ 0.12. Several important conclusions may be drawn:  \\bigskip  {\\bf (1)} For every combination, the $PL$-relation in $K$ gives the shortest distance, and the difference between the distance based on $V$ and $K$ can be as large as 0.24 mag (solutions 4, 5 in Table~8).   This systematic effect is worrying and merits investigation. It could hint to errors in the reddening, or dereddening procedure. It is illustrative to note that the (minimum, maximum, mean) extinction for the whole sample is ($-0.21$, 3.4, 1.3) in $A_{\\rm V}$, ($-0.08$, 2.1, 0.75) in $A_{\\rm I}$, and ($-0.05$, 0.26, 0.10) in $A_{\\rm K}$. This implies that any uncertainty in reddening is less in $K$. In Sect.~5.1 we have estimated these uncertainties (about 0.04 in $V$, 0.03 in $I$ and 0.01 mag in $K$) and added them as a random errors. Possibly these are errors of a systematic nature instead. It is interesting to note that application of the procedures outlined in Sect.~3.2 results in negative reddening in some cases.  One can raise the question how much bluer the period-colour-relations need to be to give positive reddening for all stars. It turns out that Eq.~(4) needs to be bluer by 0.065 mag, Eq.~(8) by 0.055 mag, and Eq.~(11) by 0.075 mag. For the default slopes and using all stars, the zero points in $V,I,K$ would change to, respectively, $-1.620$ (from $-1.411$), $-1.970$ (from $-1.892$), and $-2.660$ (from $-2.607$).  On the other hand, the DM based on the Wesenheit-index, which avoids the use of a $PC$-relation to estimate the individual reddenings, is in perfect agreement with the DM based both on the $PL$-relations in $V$ and $I$. This suggests that the reddening is not the main reason for the systematically shorter DM in the $K$-band.  As pointed out in Sect.~4.3 there may be a bias in the $K$-band zero point due to the smaller number of \\C\\ with accurate intensity-mean magnitudes. The correction for this bias was estimated by comparing, for the same sample of stars with $K$-band data, the zero point of the $PL$-relations in $V$ and $I$ to those for the full samples in $V$ and $I$, and this correction is about 0.1 mag, in the sense that it makes the DM based on $K$ shorter than they would be without this correction. This uncertainty can only be eliminated if more intensity-mean NIR magnitudes come available.      {\\bf (2)} The uncertainty in the type of metallicity correction introduces a range in DM of up to 0.20 mag in $V,I$, and 0.12 mag in $K$.  {\\bf (3)} The uncertainty in the slope of the Galactic $PL$-relations introduces a range in DM of about 0.16 mag in $V,I$, and about 0.05 mag in $K$.\\\\   \\noindent Taking the case with the observed slopes of LMC Cepheids with no metallicity correction as default, one may summarise the results as follows.  Based on the $PL$-relation in $V$ and $I$, and the Wesenheit-index, the true DM to the LMC is 18.60 $\\pm$ 0.11 ($\\pm$ 0.08 slope) ($^{+0.08}_{-0.15}$ metallicity). Based on the $PL$-relation in $K$ it is 18.52 $\\pm$ 0.18 ($\\pm$ 0.03 slope) ($\\pm 0.06$ metallicity) ($^{+0.10}_{-0}$ sample bias). The terms between parenthesis indicate the possible systematic uncertainties due to the slope of the Galactic $PL$-relations, the metallicity corrections, and in the $K$-band, due to the limited number of stars. Recent work by Sandage et al. (1999) indicate that the effect of metallicity towards shorter distances may be smaller in $V$ and $I$ than indicated here. A more accurate determination is not possible without more definite information on the slope of the Galactic $PL$-relations and the metallicity correction. \\\\  Our prefered distance modulus is the one based on the $PL$-relation in $V$, $I$ and the Wesenheit index, and puts the LMC 0.10 mag in DM closer than the value of 18.70 derived by FC. The difference is due to four effects that all work in the same direction, namely, (1) FC apply a metallicty correction of +0.042 mag, (2) the difference in the zero point in the Galactic $PL$-relation in $V$ between FC and this study is +0.03 mag (+0.01 mag is due to Malmquist bias which FC did not take into account, while +0.02 mag is due to the different sample and slightly different photometry in some cases), (3) the difference in the DM based on $V$ compared to the mean of the DM based on $V$, $I$ and the Wesenheit-index is +0.02 mag, and (4) the difference between FC and this study in the adopted zero point of the $PL$-relation in $V$ of the LMC Cepheids is +0.01 mag. \\\\  We finally note that the influence of the choice of slope and the metallicity correction are (predicted to be) smallest in the $K$-band as well as the uncertainty in the extinction correction. If the 20-30 closest \\C\\ without published NIR photometry could have their NIR intensity-mean magnitudes determined, then the uncertainty due to the small number of stars could be eliminated and the zero point in $K$ could be determined with an error that is a factor of two smaller than is possible at present. \\\\   \\begin{table} \\caption[]{Metallicity dependence of the absolute magnitude between Galaxy and LMC from recent theoretical models. } \\begin{tabular}{rrrrl} \\hline $\\log P_0$ & ${\\Delta M}_{\\rm V}$ & ${\\Delta M}_{\\rm I}$ & ${\\Delta M}_{\\rm K}$ & Reference \\\\ \\hline 0.77 &  0.139 & 0.145 & 0.073 & Bono et al. (1999)\\\\ 0.77 &  0.005 &$-0.007$ & 0.004 & Alibert et al. (1999)\\\\ 0.47 & $-0.032$ & 0.062 & 0.025 & Bono et al. (1999)\\\\ 0.47 & $-0.008$ &$-0.018$ &$-0.003$ & Alibert et al. (1999)\\\\ \\hline \\end{tabular} \\end{table}   \\begin{table*} \\caption[]{Distance Moduli to the LMC. Based on the $PL$-relations in $V,I,K$ and the Wesenheit-index, for different assumptions about the slope of the Galactic $PL$-relation, and metallicity correction.} \\begin{tabular}{crccccccc} \\hline Solution & Slope & Metallicity & $V$ & $I$ & $K$ & $W$ & Mean\\\\ &       & dependence  &     &     &     &   & over $V,I,K$  \\\\ \\hline 1& traditional & 0 & 18.62 $\\pm$ 0.11 & 18.60 $\\pm$ 0.12 & 18.52 $\\pm$ 0.18 & 18.60 $\\pm$ 0.11 & 18.60 $\\pm$ 0.07 & \\\\ 2& traditional & + & 18.71 $\\pm$ 0.12 & 18.70 $\\pm$ 0.13 & 18.57 $\\pm$ 0.18 &         & 18.68 $\\pm$ 0.08 & \\\\ 3& traditional & $-$ & 18.47 $\\pm$ 0.12 & 18.45 $\\pm$ 0.12 & 18.45 $\\pm$ 0.18 &         & 18.46 $\\pm$ 0.08 & \\\\ 4& shallower & 0 & 18.80 $\\pm$ 0.14 & 18.76 $\\pm$ 0.19 & 18.57 $\\pm$ 0.19 &         & 18.73 $\\pm$ 0.10 & \\\\ 5& shallower & + & 18.86 $\\pm$ 0.12 & 18.86 $\\pm$ 0.16 & 18.62 $\\pm$ 0.19 &         & 18.81 $\\pm$ 0.09 & \\\\ 6& shallower & $-$ & 18.65 $\\pm$ 0.14 & 18.61 $\\pm$ 0.18 & 18.50 $\\pm$ 0.19 &        & 18.60 $\\pm$ 0.10 & \\\\ 7& steeper   & 0 & 18.66 $\\pm$ 0.13 & 18.64 $\\pm$ 0.15 & 18.53 $\\pm$ 0.19 &        & 18.63 $\\pm$ 0.09 & \\\\ 8& steeper   & + & 18.72 $\\pm$ 0.18 & 18.75 $\\pm$ 0.16 & 18.58 $\\pm$ 0.20 &        & 18.70 $\\pm$ 0.10 & \\\\ 9& steeper   & $-$ & 18.51 $\\pm$ 0.14 & 18.49 $\\pm$ 0.15 & 18.46 $\\pm$ 0.19 &        & 18.49 $\\pm$ 0.09 & \\\\ \\hline mean/range & & & 18.67 / 0.20 & 18.65 / 0.21 & 18.53 / 0.09 & & 18.63 \\\\ \\hline \\end{tabular} \\end{table*}     \\subsection*{Acknowledgements} We thank Pascal Fouqu\\'e, Michael Feast and Patricia Whitelock for providing tabular material in Gieren et al. (1998), respectively Feast \\& Whitelock (1997) in electronic format.  We thank Giuseppe Bono and Santi Cassisi in calculating and communicating additional $PL(C)$-relations to us.  Frederic Pont and Frederic Arenou are thanked for lively and interesting discussions.  This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France."
    },
    "0002/astro-ph0002055_arXiv.txt": {
        "abstract": "We consider gravitational waves emitted by various populations of compact binaries at cosmological distances. We use population synthesis models to characterize the properties of double neutron stars, double black holes and double white dwarf binaries as well as white dwarf-neutron star, white dwarf-black hole and black hole-neutron star systems. \\\\ We use the observationally determined cosmic star formation history to reconstruct the redshift distribution of these sources and their merging rate evolution.\\\\ The gravitational signals emitted by each source during its early-inspiral phase add randomly to produce a stochastic background in the low frequency band with spectral strain amplitude between $\\sim 10^{-18} \\, \\mbox{Hz}^{-1/2}$ and $\\sim 5 \\times 10^{-17}\\,\\mbox{Hz}^{-1/2}$ at frequencies in the interval $[\\sim 5 \\times 10^{-6}-5 \\times 10^{-5}]$~Hz.  The overall signal which, at frequencies above $10^{-4}$~Hz, is largely dominated by double white dwarf systems, might be detectable with LISA in the frequency range $[1-10]$~mHz and acts like a confusion limited noise component which might limit the LISA sensitivity at frequencies above 1~mHz. ",
        "introduction": "Binaries with two compact stars are the most promising sources for gravitational radiation. The final phase of spiral in may be detected with ground-based (LIGO, VIRGO, GEO and TAMA) and space-borne laser interferometers (LISA). This has motivated researchers to model gravitational waveforms and to develop population synthesis codes to estimate the properties and formation rates of possible sources for gravitational wave radiation.  Since there is not yet a single prescription for calculating the gravitational emission from a compact binary system, it is customary to divide the gravitational waveforms in two main pieces: the inspiral waveform, emitted before tidal distortions become noticeable, and the coalescence waveform, emitted at higher frequencies during the epoch of distortion, tidal disruption and/or merger (Cutler \\etal 1993).  As the binary, driven by gravitational radiation reaction, spirals in, the frequency of the emitted wave increases until the last 3 orbital cycles prior to complete merger.  With post-Newtonian expansions of the equations of motion for two point masses, the waveforms can be computed fairly accurately in the relatively long phase of spiral in (see, for a recent review, Rasio \\& Shapiro 2000 and references therein). Conversely, the gravitational waveform from the coalescence can only be obtained from extensive numerical calculations with a fully general relativistic treatment. Such calculations are now well underway (Brady, Creighton \\& Thorne 1998; Damour, Iyer \\& Sathyaprakash 1998; Rasio \\& Shapiro 1999).  In this paper, we consider the low frequency signal from the early phase of the spiral in, which is of interest for space-borne antennas, such as LISA. For this purpose, we use the leading order expression for the single source emission spectrum, obtained using the quadrupole approximation. Our analysis includes various populations of compact binary systems: black hole-black hole (bh, bh), neutron star-neutron star (ns, ns), white dwarf-white dwarf (wd, wd) as well as mixed systems such as (ns, wd), (bh, wd) and (bh, ns).  For some of these sources [(ns, ns), (wd, wd) and (ns, wd)], statistical information on the event rate can be inferred from electromagnetic observations.  In particular, there are several observational estimates of the (ns, ns) merger rate obtained from statistics of the known population of binary radio pulsars (Narayan, Piran \\& Shemi 1991; Phinney 1991).  A rather large number of close white dwarf binaries have recently been found (see Maxted \\& Marsh 1999 and Moran 1999). However, it is customary to constrain the (wd, wd) merger rate from the observed SNIa rate (see Postnov \\& Prokhorov 1998).  Also the population of binaries where a radio pulsar is accompanied by a massive unseen white dwarf may be considerably higher than hitherto expected (Portegies Zwart \\& Yungelson 1999).    Since most stars are members of binaries and the formation rate of potential sources of gravitational waves may be abundant in the Galaxy, the gravitational-wave signal emitted by such binaries might produce a stochastic background.  This possibility has been explored by various authors, starting from the earliest work of Mironovskij (1965) and Rosi \\& Zimmermann (1976) until the more recent investigations of Hils, Bender \\& Webbink (1990), Lipunov \\etal (1995), Bender \\& Hils (1997), Giazotto, Bonazzola \\& Gourgoulhon (1997), Giampieri (1997), Postnov \\& Prokhorov (1998),  and Nelemans, Portegies Zwart \\& Verbunt (1999).  This background, which acts like a noise component for the interferometric detectors, has always been viewed as a potential obstacle for the detection of gravitational wave backgrounds coming from the early Universe.  In this paper we extend the investigation of compact binary systems to extragalactic distances, accounting for the binaries which have been formed since the onset of galaxy formation in the Universe.  Following Ferrari, Matarrese \\& Schneider (1999a, 1999b: hereafter referred as FMSI and FMSII, respectively), we modulate the binary formation rate in the Universe with the cosmic star formation history derived from observations of field galaxies out to redshift $z \\sim 5$ (see \\eg Madau, Pozzetti \\& Dickinson 1998b; Steidel \\etal 1999).  The magnitude and frequency distribution of the integrated gravitational signal produced by the cosmological population of compact binaries is calculated from the distribution of binary parameters (masses and types of both stars, orbital separations and eccentricities). These orbital parameters characterize the gravitational-wave signal which we observe on Earth.  Detailed information for the properties of the binary population may be obtained through population synthesis.  We use the binary population synthesis code {\\sf SeBa} to simulate the characteristics of the binary population in the Galaxy (Portegies Zwart \\& Verbunt 1996; Portegies Zwart \\& Yungelson 1998). The characteristics of the extragalactic population are derived from extrapolating these results to the local Universe.   The paper is organized as follows: in Section~2 we describe the population synthesis calculations. Section~3 deals with the energy spectrum of a single source followed, in Section~4, by the derivation of the extragalactic backgrounds for the different binary populations. In Sections~3 and~4 we also give details on the adopted astrophysical and cosmological properties of the systems. In Section~5, we compute the spectral strain amplitude produced by each cosmological population and investigate its detectability with LISA.  Finally, in Section~6 we summarize our main results and compare them with other previously estimated astrophysical background signals. ",
        "conclusions": "In this paper we have obtained estimates for the stochastic background of gravitational waves emitted by cosmological populations of compact binary systems during their early-inspiral phase.  Since we have restricted our investigation to frequencies well below the frequency emitted when each system approaches its last stable circular orbit, we have characterized the single source emission using the quadrupole approximation.  Our main motivation was to develop a simple method to estimate the gravitational signal produced by populations of binary systems at extragalactic distances. This method relies on three main pieces of information: \\begin{enumerate} \\item the theoretical description of gravitational waveforms to characterize the single source contribution to the overall background \\item the predictions of binary population synthesis codes to characterize the distribution of astrophysical parameters (masses of the stellar components, orbital parameters, merger times etc.) among each ensemble of binary systems \\item a model for the evolution of the cosmic star formation history derived from a collection of observations out to $z \\sim 5$ to infer the evolution of the birth and merger rates for each binary population throughout the Universe. \\end{enumerate} As we have considered only the early-inspiral phase of the binary evolution, our predictions for the resulting gravitational signals are restricted to the low frequency band $10^{-5}-1$~Hz. The stochastic background signals produced by (wd, wd) and (ns, ns) might be observable with LISA and add as confusion limited noise components to the LISA instrumental noise and to the signal produced by binaries within our own Galaxy. The extragalactic contributions are dominant at frequencies in the range $1-10$~mHz and limit the performances expected for LISA in the same range, where the previously estimated sensitivity curve was attaining its mimimum.  We plan to extend further this preliminary study and to consider more realistic waveforms so as to enter a frequency region interesting for ground-based experiments.  Finally, in Fig.~\\ref{fig:extragal} we show the spectral densities of the extragalactic backgrounds that have been investigated so far. The high frequency band appears to be dominated by the stochastic signal from a population of rapidly rotating neutron stars via the r-mode instability (see FMSII). For comparison, we have shown the overall signal emitted during the core-collapse of massive stars to black holes (see FMSI). In this case, the amplitude and frequency range depend sensitively on the fraction of progenitor star which participates to the collapse. The signal indicated with BH corresponds to the conservative assumption that the core mass is $\\sim 10 \\%$ of the progenitor's (see FMSI). Recent numerical simulations of core-collapse supernova explosions (Fryer 1999) appear to indicate that for progenitor masses $>40 \\msun$ no supernova explosion occurs and the star directly collapses to form a black hole. The final mass of this core depends strongly on the relevance of mass loss caused by stellar winds (Fryer \\& Kalogera 2000). If massive black holes are formed the resulting background would have a larger amplitude and the relevant signal would be shifted towards lower frequencies, more interesting for ground-based interferometers (Schneider, Ferrari \\& Matarrese 1999).  In the low frequency band, we have plotted only the backgrounds produced by (bh, bh), (ns, ns) and (wd, wd) binaries because their signals largely overwhelm those from other degenerate binary types. \\begin{figure} \\centerline{\\psfig{figure=extragalactic.ps,angle=270,width=8cm}} \\caption{The predicted strain amplitude of the stochastic backgrounds produced by extragalactic populations of gravitational sources. In the high frequency band, we show the estimates for the background produced by rotating neutron stars via r-mode instability, and two possible signals emitted by populations of massive stars collapsing to black holes (see text). In the low frequency band, we plot the background predicted for three different populations of binary systems.} \\label{fig:extragal} \\end{figure}  We find that both in the low and in the high frequency band, extragalactic populations generate a signal which is comparable to and, in some cases, larger than the backgrounds produced by populations of sources within our Galaxy (Giazotto, Bonazzola \\& Gourgoulhon 1997; Giampieri 1997; Postnov 1997; Hils, Bender \\& Webbink 1990; Bender \\& Hils 1997; Postnov \\& Prokhorov 1998; Nelemans, Portegies Zwart \\& Verbunt 1999). It is important to stress that even if future investigations reveal that the amplitude of galactic backgrounds might be higher than presently conceived, their signal could still be discriminated from that generated by sources at extragalactic distance. In fact, the signal produced within the Galaxy shows a characteristic amplitude modulation when the antenna changes its orientation with respect to fixed stars (Giazotto, Bonazzola \\& Gourgoulhon 1997; Giampieri 1997).   The same conclusions can be drawn when the extragalactic backgrounds are compared to the stochastic relic gravitational signals predicted by some classical early Universe scenarios. The relic gravitational backgrounds suffer of the many uncertainties which characterize our present knowledge of the early Universe. According to the presently conceived typical spectra, we find that their detectability might be severely limited by the amplitude of the more recent astrophysical backgrounds, especially in the high frequency band."
    },
    "0002/astro-ph0002263_arXiv.txt": {
        "abstract": "{Galaxy interactions, mergers, elliptical formation, bulge formation, starbursts, quasars} Gravitational interactions and mergers are shaping and reshaping galaxies throughout the observable universe.  While observations of interacting galaxies at low redshifts yield detailed information about the processes at work, observations at high redshifts suggest that interactions and mergers were much more frequent in the past.  Major mergers of nearby disk galaxies form remnants that share many properties with ellipticals and are, in essence, present-day protoellipticals.  There is also tantalizing evidence that minor mergers of companions may help build bulges in disk galaxies.  Gas plays a crucial role in such interactions. Because of its dissipative nature, it tends to get crunched into molecular form, turning into fuel for starbursts and active nuclei.  Besides the evidence for ongoing interactions, signatures of past interactions and mergers in galaxies abound: tidal tails and ripples, counterrotating disks and bulges, polar rings, systems of young globular clusters, and aging starbursts.  Galaxy formation and transformation clearly is a prolonged process occurring to the present time.  Overall, the currently available observational evidence points towards Hubble's morphological sequence being mainly a sequence of decreasing merger damage. ",
        "introduction": "Ever since Hubble (1936) published his famous `Sequence of Nebular Types' (a.k.a.\\ tuning-fork diagram) the question has been: What determines the position of a galaxy along this sequence?  And why are galaxies at one end of the sequence disk-shaped and at the other end ellipsoidal?  Was this shape dichotomy imprinted during an early collapse phase of galaxies, or did it arise through subsequent evolution?  Work begun several decades ago by Zwicky (1956), Arp (1966), Alladin (1965), and Toomre \\& Toomre (1972, hereafter `TT'), among others, has led to growing evidence that gravitational interactions between neighbor galaxies do not only explain some of the most striking `bridges' and `tails' observed in disturbed galaxy pairs, but also tend to lead to galactic mergers that often trigger bursts of star formation and clearly represent important phases of galaxy building (Larson 1990; Barnes \\& Hernquist 1992; Kennicutt \\etal\\ 1998).  Before reviewing some of the evidence for interactions and mergers being a significant driver of galaxy evolution, it seems wise to agree on some terminology and point out biases.  To be called a `merger', a galaxy pair or single galaxy should show at least clear morphological signatures of an advanced tidal interaction, such as significant distortions, major tails, and ripples or `shells' (for a review, see Schweizer 1998).  A stronger case for merging can usually be made when {\\it kinematic\\,} signatures are available as well, such as opposite tail motions, counter-rotating parts, or tail material falling back onto a remnant.  As figure~\\ref{fig:n4038kin} illustrates, much recent progress in this area is due to the upgraded {\\it Very Large Array}'s ability to map the line-of-sight motions of neutral hydrogen (\\hi ) in tidal features in great detail (Hibbard 1999).  \\begin{figure} \\centerline{\\psfig{file=schweizer_fig1.ps,height=7.0cm,angle=-90}} \\caption{Neutral hydrogen distribution and kinematics of NGC\\,4038/4039. Left: \\hi\\ contour lines ({\\it white}) superposed on an optical photograph; right: \\hi\\ position--velocity plot, with declination along $y$-axis and line-of-sight velocity along $x$-axis (from Hibbard 1999).} \\label{fig:n4038kin} \\end{figure}  The main bias in studies of gravitational interactions has been toward major mergers, which involve two galaxies of nearly equal mass.  Such mergers are highly destructive and tend to lead to spectacular morphologies, whence they can be observed from the local universe out to redshifts of $z\\approx 2$ and beyond.  Minor mergers involving galaxies with mass ratios of, say, $m/M = 0.1$--0.5 are less spectacular and often require verification via some kinematic signature (esp.\\ in the remnant phase).  Hence, such interactions and mergers have been studied mainly in nearby galaxies and out to $z\\lesssim 0.5$. Finally, although satellite accretions leading to mass increases of a few percent or less may be relatively frequent, they are the most difficult to detect and have been studied only in the Local Group, and even there nearly exclusively in our Milky-Way galaxy.  Thus, our knowledge of growth through accretions and minor mergers is severely limited.  Because of its dissipative nature gas plays a disproportionately large role in galaxy interactions.  Even at the present epoch, the vast majority of galaxies contain significant amounts of cold gas (Roberts \\& Haynes 1994).  During tidal interactions and mergers this gas tends to be driven toward the centers of galaxies through gravitational torques exerted on it by tidally induced {\\it stellar\\,} bars (e.g. Barnes \\& Hernquist 1996). The ensuing shocks and energy dissipation allow the gas to get compressed, leading to intense bursts of star formation, globular-cluster formation, and the feeding of nuclear activity.  Starbursts and active galactic nuclei in turn drive galactic winds and jets, which can have profound effects on the chemical evolution of galaxies (Heckman 2000).  Some of these processes can now be reproduced by modern $N$-body simulations that include gas hydrodynamics.  Barnes (1999) shows a beautiful sequence of two gas-rich disk galaxies merging.  Whereas their stars end up in a three-dimensional pile not unlike an elliptical galaxy with considerable fine structure, more than half of the cold gas from the input disks gets funneled to the center of the remnant into a region only about 0.5\\3kpc in diameter, while the initially warm gas ($T\\approx 10^4$\\3K) gets heated to X-ray temperatures ($\\sim$10$^6$\\3K) and forms a pressure-supported atmosphere of similar dimensions as the stellar pile. The time scale for this transformation from two disk galaxies to one merged remnant is remarkably short: about 1.5 rotation periods of the input disks or, when scaled to component galaxies of Milky-Way size, about 400\\3Myr.  The rapidity of this equal-mass merger is due to strong dynamical friction. We should keep this in mind when trying to understand the formation of elliptical galaxies in dense environments.  Claims have been made that cluster ellipticals formed in a rapid monolithic collapse because their present-day colors are rather uniform.  Yet, experts agree that age differences of $\\lesssim$3\\3Gyr cannot be discerned from broad-band colors of galaxies 10--15\\3Gyr old.  A time interval of 3\\3Gyr may seem short when we struggle with logarithmic age estimates, yet it is long when compared to the merger time scale.  About eight major mergers of the kind simulated by Barnes could take place one after another during this time interval, and 12\\3Gyr later all their remnants would look nearly the same color and age.  Hence, claims about monolithic collapses and a single epoch of elliptical formation are to be taken with a grain of salt.  There was time for many major mergers of juvenile disks during the first few Gyr after the big bang, and most cluster ellipticals could have formed through such mergers without us knowing it from their present-day colors.  The following review of evidence for interactions being a driver of galaxy evolution begins with accretions in the Local Group, continues with minor mergers and the damage they inflict on disk galaxies, moves on to major mergers forming ellipticals from wrecked disks, and ends with a brief description of what we have learned from first glimpses of high-redshift mergers. ",
        "conclusions": "This review has high-lighted the role that interactions and mergers play in driving galaxy evolution.  At present we remain challenged to understand the relative importance of weak and strong interactions, the details of bulge formation, the existence of nearly pure-disk galaxies, and the merger rate as a function of redshift.  Yet, some firm conclusions have been reached and are as follows:  \\begin{itemize} \\item Gravitational interactions and mergers are forming and transforming galaxies throughout the observable universe.  The vast majority involve gas, dissipation, and enhanced star formation.  \\item The close link between mergers, ultra-luminous infrared galaxies, and quasars suggests that---like quasar activity---major merging may have peaked around $z\\approx 2$.  \\item Major disk--disk mergers form elliptical galaxies with kinematic subsystems, bimodal globular-cluster populations, and remnant fine structure. Such mergers occurred relatively early near the centers of rich clusters, but continue to the present time in rich-cluster outskirts, poorer clusters, and the field.  \\item Minor mergers tend to move disk galaxies toward earlier morphological types, creating kinematic subsystems and some bulges (fraction remains unknown).  \\item In short, the currently available evidence strongly suggests that Hubble's morphological sequence is mainly a sequence of decreasing merger damage.  \\end{itemize}"
    },
    "0002/astro-ph0002505_arXiv.txt": {
        "abstract": "Accreting  black holes and neutron  stars in their hard (low) state show not only very similar X/$\\gamma$-ray spectra but also that the behaviour of their light curves is quite similar which can be quantified as having similar power-density spectra and Fourier-frequency-dependent time/phase lags.  Taken  together  this  argues  for  a  common  mechanism  of  the X/$\\gamma$-ray  production in these objects.  This mechanism is probably a property  of the  accretion  flow only since it does not depend on the nature of the compact object.  In this paper, I review the observational data paying most attention to the properties of the temporal variability such as the time/phase  lags that hopefully can help us to  discriminate between  different   theoretical  models.  I  also  discuss  the  models developed to account for the basic observational facts.  Particularly, I show  that  the  commonly  used  Compton  cloud  models  with   constant temperature cannot explain variable sources without violating the energy conservation law.  Alternative models where time lags are related to the spectral  evolution  during X-ray flares are discussed and compared with observations.  Compton  reflection  from the outer edge of the accretion disc is shown to markedly affect the time lag Fourier spectrum. ",
        "introduction": "\\label{sec:intro}  X-ray and gamma-ray  spectra of accreting  black holes and neutron stars are  deconvolved  into  (at  least)  two  components:  a soft  component interpreted as emission from an optically  thick  accretion  disc, and a hard tail  associated  with a hot (10--100 keV)  ``corona''.  Reviews of the spectral  properties of Galactic black hole candidates (GBHs) can be found in Gilfanov et al.  (1995),  Tanaka \\& Lewin  (1995),  Grebenev et al.  (1993,   1997),   Grove  et  al.  (1997),  and   Poutanen   (1998). X/$\\gamma$-ray  properties of radio-quiet  active galactic  nuclei (AGN) are reviewed by  Zdziarski  et al.  (1997),  Johnson et al.  (1997), and Zdziarski (1999).  Recent results on the broad-band spectra of accreting neutron stars are presented by Barret et al.  (2000).  An  amusing  fact is that  super-massive  black  holes in AGN,  GBHs and accreting  neutron stars in their hard (low) states (see Tanaka \\& Lewin 1995; Gilfanov et al.  1995 for the  definition of the spectral  states) show very similar X/$\\gamma$-ray  spectra (see Zdziarski 1999; Barret et al.  2000).  Furthermore, properties of their rapid temporal variability are also  similar  (van der Klis 1995b;  Wijnands  \\& van der Klis 1999; Psaltis,  Belloni \\& van der Klis  1999;  Ford et al.  1999;  Edelson \\& Nandra  1999;  Chiang  et  al.  2000).  All  this  argues  for a  common mechanism of the X-ray production in all these sources.    There are good reasons to believe that the main radiative  mechanism for the  production  of the hard X-rays is  Comptonization  of soft  photons (e.g.,  Shapiro,  Lightman, \\& Eardley 1976; Sunyaev \\& Tr\\\"umper  1979; Sunyaev \\& Titarchuk  1980).  However, it is not  completely  clear what determines  the  observed  spectral  slopes.  The  geometry of the X-ray emitting  region  and the  source of soft  photons  is still a matter of debate (see Svensson 1996; Poutanen 1998; Beloborodov  1999b; Wardzinski \\& Zdziarski 2000).   An important clue to our understanding of the X-ray production came from the  discovery of Fe lines (at $\\sim 6.4$ keV) and the  hardening of the spectra  above 10 keV in AGN (Pounds et al.  1990;  Mushotzky,  Done, \\& Pounds  1993;  Nandra \\& Pounds  1994),  Cygnus X-1  (e.g.,  Done et al. 1992;  Gierli\\'nski  et al.  1997), and neutron stars (e.g.,  Yoshida et al.  1993).  These features are associated  with the  reflection of hard X-rays from cold material (Basko,  Sunyaev, \\& Titarchuk 1974; George \\& Fabian  1991;  Magdziarz  \\&  Zdziarski  1995;  Poutanen,  Nagendra,  \\& Svensson  1996).  These  observations  gave  support  to the  so  called two-phase  accretion  disc-corona  models.  In such  models,  X-rays are emitted by a hot rarified  corona above the cold accretion  disc (Haardt \\& Maraschi  1993; Haardt,  Maraschi, \\&  Ghisellini  1994; Stern et al. 1995;  Poutanen \\& Svensson  1996).  Hard X-rays from the corona,  being reprocessed  in the cold disc,  produce the  reflection  hump as well as most of the seed  soft  photons  that are  subsequently  Comptonized  to produce  the hard  X-rays.  This is the {\\it  feedback}  mechanism.  The geometry of the corona  determines the feedback factor which in its turn determines   the  spectral   slope  of  the   escaping   radiation.  The temperature of the emitting  plasma (or to be more exact, the Kompaneets $y$-parameter)  is determined by the energy balance between  heating (by magnetic   reconnection?)  and  cooling  (by   Comptonization   of  soft photons).  Further support for the feedback models was recently given by Zdziarski, Lubinski, \\& Smith (1999) (see also Zdziarski 1999; Gilfanov,  Churazov, \\&  Revnivtsev  2000) who found a  correlation  between  the  amount  of reflection  ($R\\equiv\\Omega/(2\\pi)$,  where  $\\Omega$  is a solid  angle subtended  by cold  material as viewed  from the X-ray  source)  and the intrinsic photon spectral index, $\\Gamma$, of the hard X-ray  component. Such a correlation  can easily be explained if there is overlap  between the hot corona and the cold disc (Poutanen,  Krolik, \\& Ryde 1997).  The further  the cold disc  penetrates  into the  corona,  the larger is the cooling, the smaller is the temperature of the corona, the softer is the spectrum,  and, finally, the larger is the amplitude of the  reflection. The model, however, appears to have trouble giving reflection amplitudes above  $R_{\\max}\\sim  0.5$ (if the coronal optical depth $\\taut \\sim 1$, see  Zdziarski et al.  1997) due to partial  smearing of the  reflection component by the hot corona.    Alternatively,  the observed $R - \\Gamma$  correlation can be reproduced by variations of the bulk velocity of the X/$\\gamma$-ray emitting plasma (Beloborodov 1999a,b).  If the emitting regions are sufficiently compact to  produce  electron-positron  pairs,  the  pressure  of the  radiation reflected  and  reprocessed  in the disc  accelerates  pairs  to  mildly relativistic  velocities  away from the  disc.  On the  other  hand, for proton dominated  plasmas, a small anisotropy in the energy  dissipation mechanism  can result in the ejection of  particles  away or towards the disc.  Ejection away from the disc reduces $R$ below 1 and leads to hard spectra, while ejection towards to the disc can result in apparent $R>1$ as is observed in some objects.      The physical  possibility of the corona formation was studied by Galeev, Rosner, \\& Vaiana  (1979).  They showed that the magnetic  fields, being amplified  in the cold disc due to  turbulent  motions and  differential rotation, do not have time to  annihilate  inside the disc on the inflow time  scale.  Instead,  the field  loops are  expelled  from the disc by buoyancy (the Parker  instability)  and they  annihilate  in the tenuous corona.  Beloborodov (1999a) showed that the mechanism studied by Galeev et al.  is able to produce a corona of limited luminosity which is $h/r$ (the ratio of the disc height to its radius) times smaller than the disc luminosity.  By contrast, in some sources most of the energy  escapes in the   form  of  hard   X-rays.  Beloborodov   also   argued   that   the magneto-rotational   instability  (Velikhov  1959;  Chandrasekhar  1960; Balbus  \\&  Hawley  1991)  increases  the  rate  of the  magnetic  field generation (as compared with the Galeev et al.  model) thus producing an active  magnetic  corona  where a large  fraction  of the  gravitational energy can finally be dissipated in magnetic flares.  These  qualitative arguments  were  recently   supported  by  numerical   three-dimensional magnetohydrodynamical  simulations  of Miller \\& Stone (2000) who showed that  about 25 \\% of the  total  energy  dissipation  can  occur  in the rarified corona.    An alternative to the magnetic  corona is the hot disc model (Shapiro et al.  1976;  Ichimaru  1977;  Narayan,   Mahadevan,  \\&  Quataert   1998; Zdziarski  1998; Esin et al.  1998)  which is also able to  explain  the observed  X/$\\gamma$-ray  spectra.  In order to distinguish  between the models, it would be helpful to  compare  the  predictions  of  different models with the temporal  variability data (see van der Klis 1995a,b and Cui 1999a for recent reviews).  Unfortunately, most of the papers on the spectral models do not consider the temporal  variability.  On the other hand, most of the models  designed to explain the  temporal  variability data do not pay  enough  attention  to the  emission  processes  and the physics of the spectral formation.   In this review, we will  discuss the  variability  data  keeping in mind recent  advances  in  modelling  broad-band  X/$\\gamma$-ray  spectra  of accreting black holes and neutron stars.  Most attention will be paid to the time  lags  that  can  shed  light  on the  mechanism  of the  X-ray production.  Then we discuss  simple  phenomenological  models  that are able to explain some of the observational  facts.  After that, we switch to  the  physical   models.  In  particular,   the   properties  of  the Comptonizing  regions will be discussed.  We will point out the flaws in models that do not consider the energy balance in the ``Compton cloud'', and then discuss  models that  satisfy the energy  conservation  law and confront them with the available data. ",
        "conclusions": "Time lags and other temporal variability data provide strong constraints on the models of the X-ray production.  It was demonstrated  that static Compton cloud models are based on  physically  unrealistic  assumptions. The models  invoking  spectral  evolution of the flare  spectrum can fit both the CCF and the time lag  Fourier  spectra  only if (1) the  energy dissipation  rate  increases  slowly and  decreases  rapidly and (2) the flare  spectrum  evolves  from soft to hard.  If soft seed  photons  are produced by  reprocessing  the hard ones, the change of sign in the time lag spectrum is expected at high frequencies  corresponding to the light crossing  time of the  emission  region.  The  absence  of such a change would put constraints on the size of the emitting region.  We also argued that the  reflection  of hard X-rays from the outer part of the accretion  disc produces  time delays that we already  might have observed in GBHs.  If so, the disc should be flared and the break in the time lag Fourier  spectra then  corresponds to the size of the accretion disc.  Of course, such an interpretation  is not unique.  Alternatively, small scale spectral transitions (e.g., oscillations of the inner radius of the accretion  disc at viscous time scales)  might  produce time lags observed at lower frequencies.  In the case of (quasi-)  periodic  oscillations  from the  neutron  star sources, we argued that in order to reproduce both the time lags and the energy dependent rms amplitude, the spectrum of the hot spots should not be close to a black-body."
    },
    "0002/astro-ph0002169_arXiv.txt": {
        "abstract": "We present BeppoSAX observations of the Seyfert 1.8 galaxy NGC1365 in the 0.1--100 keV range. The source was 6 times brighter than during an ASCA observation 3 years earlier. The 4--10 keV flux is highly variable during the BeppoSAX observation, while the soft (0.1-4 keV) emission is constant within the errors. Both a cold and a warm reflector and a cold absorber are required to explain the observed spectrum. The comparison between ASCA and BeppoSAX spectra strongly suggests that the circumnuclear material has a more complex structure than a simple homogeneous torus, with quite different absorbing gas columns along different lines of sight. A broad iron K$_\\alpha$ line is also present in the spectrum, with the peak energy significantly redshifted. This can be explained by means of a relativistic disk line model. Alternatively, a warm absorption Fe line system with N$_H \\simeq 10^{23}$ cm$^{-2}$ could account for the observed line profile. ",
        "introduction": "NGC 1365 is a barred spiral galaxy (Hubble type SB0) in the Fornax cluster that hosts an active nucleus whose optical spectrum shows weak broad Balmer lines (Seyfert 1.8, Alloin et al. 1981)  In this paper we present the analysis of the spectrum of NGC 1365 in the 0.1-100 keV spectral range obtained with the BeppoSAX satellite (Boella et al. 1997).  During the past ten years NGC 1365 has been observed several times in the X-rays by ASCA (Iyomoto et al. 1997, hereafter I97) ROSAT (Komossa \\& Schulz 1998) and Ginga (Awaki 1991). The 1-10 keV continuum spectrum observed by ASCA in August 1994 and January 1995 (I97) is well reproduced by a flat powerlaw (photon index $\\Gamma$=0.8) and a thermal soft component. A strong emission feature is present at E$\\sim$6.4--7 keV, which can be fitted by a single broad emission line with E=6.58 keV and equivalent width EW=2.1 keV or, alternatively, by two narrow lines with E=6.4 keV (neutral iron, EW=0.9 keV) and E=6.7 keV (highly ionized iron, EW=0.9 keV). Both these spectral features and the lack of (short term) variability suggested that the ASCA spectrum is dominated by a cold reflection component which is usually observed in most of the heavily absorbed, Compton thick sources (Maiolino et al. 1998, hereafter M98) and generally ascribed to the reflection from the molecular torus expected by the unified model of AGNs (Antonucci 1993).  The ASCA--SIS and ROSAT--HRI data, obtained in 1994 and 1995, reveal also the presence of a strong off-nuclear X-ray source characterized by a steep powerlaw spectrum (photon index $\\Gamma$=1.7 in the 1-10 keV band) and by a strong variability on time-scales of months; during the ASCA observation in 1995 this source was as bright as the Seyfert 2 nucleus with a flux of 0.9$\\times 10^{-12}$ erg cm $^{-2}$ s$^{-1}$. The spatial resolution of the BeppoSAX instruments does not allow to separate the contribution of this source from that of the nucleus. We will discuss the possible contamination from this off-nuclear source further in Sect. 2.  In the next section we present the results of the spectral and temporal analysis of our data. In Sect. 3 we discuss the BeppoSAX data and their differences with respect to the previous X-ray observations. We assume a distance of 18.4 Mpc for NGC 1365, as estimated by Fabbiano et al. (1992), and in agreement with more recent Cepheid measurement (Madore et al. 1998). ",
        "conclusions": "We presented new BeppoSAX data in the 0.1--100 keV range of the Seyfert 1.8 galaxy NGC1365. The spectrum is characterized by a continuum absorbed by a cold gaseous column density of $\\rm N_H = 4\\times 10^{23}$ cm$^{-2}$ and an iron K$\\alpha$ emission complex that is well fitted by a cold component at 6.29 keV and a warm component at 7 keV (rest frame). At energies below the absorption cutoff (E $<$ 4 keV) a soft excess is present.  The cold absorption is probably due to the obscuring torus predicted by unified model of AGNs. The continuum is strongly variable during the BeppoSAX observation. The variability is mostly due to the hard component of the spectrum above the photoelectric cutoff (4--10 keV), while the soft component (1.65--4 keV) is essentially constant. The rapid variability very likely reflects variations of the central engine. Instead, the soft excess is probably due to an extended component, either associated to starburst activity or to hot gas in the Narrow Line Region.  The BeppoSAX spectrum is 6 times brighter than during two ASCA observations of NGC~1365 taken about 3 years earlier. The latter spectra were characterized by a flat continuum, indicative of cold Compton reflection, very likely from the circumnuclear torus.  The high reflection efficiency, deduced from the comparison of the ASCA and BeppoSAX spectra, requires a column density of the reflector much higher than that measured in absorption. We conclude that the circumnuclear medium is strongly inhomogeneous: the torus could contain Compton thick clouds or, alternatively, has a steep density gradient from the edge to the equatorial regions.  The fading of the direct emission during the ASCA observations can be explained in two ways: the central engine was hidden by a Compton thick cloud or, most probably, the nucleus was in an intrinsically low state. In the latter scenario, the temporal behavior of the cold and the warm iron lines indicate that the cold reflecting torus must be located at a distance larger than 0.15 pc, while the warm mirror must be located at a distance larger than 1 pc. Both the circumnuclear torus and the accretion disk contribute to the emission of the cold Fe line, in a proportion of about 1:2 respectively.  The cold iron line is significantly redshifted with respect to its nominal value. More specifically we measure a line peak (rest frame) of 6.29 keV, that is inconsistent with the nominal value of 6.4 keV at a significance level higher than 99\\%. A disk relativistic line can fit the observed profile, though the fit is worse than the analytical fit. Also, according to this fit the accretion disk must be oriented face on, that is an improbable geometry for an absorbed AGN like NGC 1365. Alternatively, we propose that the shift of the cold iron line is caused by a warm absorber, along the line of sight (with $\\rm N_{warm}\\approx 10^{23}$ cm$^{-2}$), that introduces an absorption Fe line at 6.5--6.7 keV: the combination of the cold emission line and the warm absorption line, convolved with the spectral resolution of BeppoSAX, results in an emission line whose center is apparently shifted at 6.29 keV. The spectral fit of the data with this second model is significantly better with respect to the relativistic disk line."
    },
    "0002/astro-ph0002443_arXiv.txt": {
        "abstract": "\\noindent We have discovered a concentration of extremely red objects (EROs; $R-K>6$) in the field of the \\mbox{$z=2.69$} quasar QSO~1213--0017 (UM~485), which is significantly overabundant compared to the field ERO surface density.  The optical/near-IR colors of the EROs and numerous other red galaxies in this field are consistent with elliptical galaxies at $z=1-2$.  \\HST\\ optical images for a subset of galaxies show regular morphologies, most of them being disky or diffuse and without any obvious evidence for interactions. Ground-based IR images show similar morphologies, indicating any dust reddening in these objects is spatially uniform. Optical spectroscopy with the W. M. Keck Telescope has found that four of the red galaxies lie at $z\\approx1.31$, and a fifth lies in the foreground at $z=1.20$. Of the $z\\approx1.31$ galaxies, one is a reddened AGN while the remaining three have rest-frame UV absorption-line spectra characteristic of old (few Gyr) stellar populations, similar to the old red galaxy LBDS~53W091 at $z=1.55$.  Including the \\mgiin\\ absorber seen in the QSO spectrum, we find five galaxies at $z\\approx1.31$ spread over $1.5~\\hfperone$~Mpc on the sky.  These results suggest we have discovered a coherent structure of old galaxies at high-redshift, possibly associated with a massive galaxy cluster. ",
        "introduction": "The observational study of galaxy formation and evolution can broadly be divided into two approaches: (1) searches for primordial galaxies, typically oriented toward high redshifts, in order to scrutinize formation and evolution processes as they occur, and (2) studies of existing old galaxies to decipher their origin and life history from their current properties.  With the advent of infrared imaging detectors, a population of infrared-bright, extremely red objects (EROs) have been uncovered which is relevant to both of these approaches.  \\citet{1988ApJ...331L..77E} discovered a few objects with exceptional optical/IR colors ($R-K\\gtrsim5$) in the first deep near-IR sky survey. While the brightest objects were shown to be $z\\leq0.8$ ellipticals \\citep{1988ApJ...332L..59L,1989ApJ...341...80E}, their tantalizing conjecture that the fainter ($K\\gtrsim18$) red objects are $z>1$ ellipticals has remained unresolved until recently.  Subsequent deep infrared imaging of primarily high-redshift radio galaxy and quasar fields serendipitously uncovered handfuls of objects with extreme colors \\citep{1992ApJ...386...52M,1992ApJ...399L..47E,1994ApJ...420L...5G, 1994AJ....107.1303H,1994ApJ...420L...1S,1995ApJ...440..515D, 1995ApJ...438L..13D}.  Deep multicolor optical/IR sky surveys have begun to assemble larger samples of these objects using well-defined selection criteria, an essential starting point for statistical studies \\citep{1994ApJ...434..114C,1998ApJ...507..558H, 1999ApJ...512...30C,1999AJ....117..102B,2000MNRAS.311..707M}. In particular, the recent availability of $1024 \\times 1024$~pixel IR detectors has made possible surveys for significant numbers of EROs \\citep{1999ApJ...523..100T}.  The use of the designation ``extremely red'' has varied in the literature. This was especially true in earlier work where the description was applied to any object appearing in IR images but undetected optically. In this paper, we follow \\citet{1996ApJ...471..720G} and use the term ``extremely red object'' (ERO) for a source with $R-K>6$.  This criteria was an operational one, encompassing several IR-detected, optically-invisible sources known at the time, and it has since become widely used.  In particular, \\citet{1999ApJ...523..100T}, who have conducted the widest deep $R-K$ survey to date, also adopt this definition. They confirm EROs defined by this criteria are unusual objects, being the reddest 2\\% of the $K\\leq20$ field galaxy population.  However, we will also pay attention in this paper to galaxies with $R-K>5$, the criteria used by \\citet{1988ApJ...331L..77E} and \\citet{1999ApJ...512...30C}, as this is about the expected color of passively-evolving elliptical galaxies at $z\\gtrsim1$.  We will generically refer to this larger sample as ``red galaxies.''  EROs are very optically faint ($R\\gtrsim24$) which has hampered studies of these objects.  Successful spectroscopy to determine their redshifts and physical nature has only become possible with the development of the Keck 10-m telescope.  Recent work has found some members of this ``missing population'' are luminous ($\\gtrsim L^*$) galaxies at $z>1$; hence determining their origin has direct relevance to the formation of massive galaxies and AGN.  Being solely defined by optical/IR color, the few EROs with measured redshifts form a heterogeneous population as expected, comprising at least two broad classes: (1) ultraluminous dusty star-forming systems, perhaps akin to local objects like Arp~220, and (2) massive galaxies with old passively-evolving stellar populations.  The best studied ERO to date, ERO~J164502+4626.4 \\citep[object 10 of][hereinafter ``HR10'']{1994AJ....107.1303H}, is the prototype for star-forming EROs.  This object is a $z=1.44$ dusty, ultraluminous galaxy with enormous ongoing star formation ($\\gtrsim1000$~\\Msun~yr\\perone) suggested by its sub-mm continuum emission \\citep{1996ApJ...471..720G,1998Natur.392..895C, 1999ApJ...519..610D}.  However, since HR~10 has among the reddest colors ($I-K=6$) of known EROs, it is unrepresentative of the bulk of the population or at least is an extreme example.  These EROs may provide excellent case studies of optically obscured star formation in the early Universe, especially given the association of faint EROs with some of the sub-mm emitting sources found by SCUBA \\citep{sma99}.  In fact, it may be that intense, very brief bursts of star formation are a common mode of star formation at high redshift.  There are also examples of EROs as galaxies with old stellar populations.  Passively evolving ellipticals at $z>1$ are expected to have large $R-K$ colors, which stem from their rising spectral energy distributions (SEDs) longward of $4000$~\\AA\\ being redshifted into the near-IR.  Therefore, selection using very red colors is an excellent method to search for early-type galaxies at $z>1$. This is especially true in searching for clusters since the surface density of galaxies becomes quite high at faint magnitudes and it would otherwise be hard to distinguish a cluster from the foreground and background populations. However, ellipticals at these distances are expected to be optically very faint so measuring spectroscopic redshifts is challenging. \\citet{dic95} has identified a cluster of elliptical galaxies associated with the powerful radio galaxy 3C~324 at $z=1.206$ along with a foreground structure at $z=1.15$.  To date, the highest-redshift collection of old, red galaxies spectroscopically confirmed has been found by \\citet{1997AJ....114.2232S}. They have discovered a $z=1.27$ cluster by its large $J-K$ colors; the cluster galaxies have $R-K\\gtrsim5$ and rest-frame UV spectra resembling local elliptical galaxies. There is also a neighboring cluster at $z=1.26$ found from its X-ray emission by \\citet{ros99} which contains spectroscopically old red galaxies.  Isolated examples of old EROs have been discovered at still higher redshifts. The very weak radio sources LBDS~53W091 at $z=1.552$ \\citep{1996Natur.381..581D,1997ApJ...484..581S} and LDBS~53W069 at $z=1.432$ \\citep{dun98,dey00} both have $R-K\\approx6$ and rest-frame UV spectra which imply ages of a few Gyr. Recently, \\citet{soi99} have identified an $R-K\\sim7$ object as an old galaxy at $z=1.58$ based on associating a large continuum break observed at $\\approx$1~\\micron\\ with redshifted 4000~\\AA\\ break.  Discovery of $z>1$ galaxies with old stellar populations offers several powerful lines of inquiry into understanding galaxy formation and evolution and its cosmological context. Detailed comparison of the absorption lines and continuum breaks of these galaxies with galaxies at $z=0$ may prove fruitful in tracing the evolutionary course and enrichment history of the oldest stellar populations.  Absolute age dating of these galaxies would provide a constraint on the time scale of galaxy formation and the age of the Universe.  Moreover, clusters of old galaxies at high redshift can provide testing grounds for competing scenarios of galaxy formation; the predicted appearance of early-type galaxies in these clusters is dramatically different in hierarchical galaxy formation scenarios as compared to monolithic collapse ones \\citep[e.g.,][]{1998MNRAS.294..705K}.  Finally, the existence of these old galaxies at high redshift potentially can constrain cosmological parameters and theories of structure formation \\citep[e.g.,][]{1998MNRAS.296.1089P}.  We are conducting an on-going study of the nature of these EROs, using deep optical and near-IR imaging from the ground to assemble a large sample of EROs for statistical study.  We have been acquiring high-resolution morphological information from \\HST\\ optical and Keck near-IR imaging, and we are obtaining spectroscopy from Keck in the optical and near-IR to determine ERO redshifts and physical properties. Deciphering the identity of EROs based on comparing broad-band colors alone to theoretical stellar population synthesis models is dubious given that the model parameter space is vast, at least comprising age, metallicity, and reddening variations. Spectroscopy is essential.  In addition, given the apparent heterogeneity of the population, a reasonably large sample of objects needs to be studied to understand the nature and relative abundances of the subsets, instead of the spectroscopy of individual EROs which has been done to date.  In this paper, we present a study of the EROs in the field of QSO~1213--0017 (RA = 12$^h$15$^m$49.8$^s$, Dec = --00\\degree 34\\arcmin 34\\arcsec\\ ; J2000.0). This $z=2.69$ quasar, also known as UM~485, has exceptionally strong and complex Mg~II absorption systems at $z=1.3196$ and $z=1.5534$ \\citep{1992ApJS...80....1S}. In \\S~2, we describe optical and near-IR imaging covering an 11~arcmin\\pertwo\\ region around this field and Keck optical spectroscopic follow-up of galaxies selected by their very red colors.  We consider in \\S~3 the surface density of the red galaxies, their morphologies, and their spectroscopic redshifts.  We examine in \\S~4 the collection of red galaxies as a whole, both their spatial distribution to consider the possibility that they are members of a cluster at $z=1.31$ and their spectrophotometric properties to understand their stellar populations.  We summarize our findings in \\S~5 and discuss their implications.  Throughout this paper, we assume a cosmology with $\\Omega=1$, $\\Lambda=0$, and $H_0=50\\ \\hf$~\\kms~Mpc\\perone. At $z=1.31$ with these parameters, 1\\arcsec\\ = 8.58~\\hfperone~kpc; the luminosity distance $d_L = 9.40$~\\hfperone~Gpc; and the angular diameter distance $d_\\theta=1.78$~\\hfperone~Gpc. ",
        "conclusions": "We have found an overdensity of EROs in the field of the $z=2.69$ quasar QSO~1213--0017 (UM~485), about a factor of 16 overdense compared to the blank field ERO surface density and at least of factor of 6 overdense at the 95\\% confidence level.  The optical/IR colors of the EROs and numerous other red galaxies in this field are consistent with those of passively-evolving elliptical galaxies at $z>1$.  \\HST\\ optical imaging shows a few of the red ($\\Rs-K>5$) galaxies seem to have early-type morphologies while the remainder are either disk galaxies or diffuse objects without any obvious core.  Their near-IR morphologies are consistent with those observed in the optical, unlike in the case of the prototypical ERO HR~10. This suggests that the dust extinction in these galaxies is either relatively smooth or that reddening by dust is not a significant cause of the colors.     Follow-up Keck spectroscopy has measured redshifts for five red galaxies.  Four lie at $z\\approx1.31$, and three of these have rest-frame UV absorption-line spectra similar to present-day elliptical galaxies, making this the most distant concentration of old galaxies spectroscopically confirmed to date. Including an \\mgiin\\ absorber seen in the spectrum of the background quasar, there are five spectroscopic redshifts at $z\\approx1.31$ with a standard deviation of 1800~\\kms\\ and a full range of 3800~\\kms\\ in the mean rest frame.  A number of lines of evidence possibly suggest the red galaxies in this field delineate the presence of a massive high-redshift cluster.  The ERO surface density is enhanced above blank field counts, and there are five spectroscopic redshifts close together.  The reddest 1213--0017 galaxies, three of which have the spectra of early-type galaxies, are nearly as luminous and as red as the brightest Coma cluster ellipticals. The red galaxies also have a large angular extent on the sky, and a considerable velocity spread, suggesting dynamical youthfulness.  An impression of youthfulness is also conveyed by the roughly filamentary distribution of the red galaxies on the sky.  Aside from the direct physical evidence, the presence of spectroscopically old EROs provides circumstantial support to the idea of a cluster.  Elliptical galaxies in the local Universe, presumably the present-day counterparts of $z>1$ old EROs, are highly clustered and predominantly found in high density regions \\citep[e.g.,][]{1980ApJ...236..351D}.  Such an analogy suggests that old EROs should be found with others of their ilk --- this idea has yet to be fully explored. Since the hallmark of rich clusters is the presence of old ellipticals, this $z=1.31$ field may be one of the most distant rich cluster candidates to date. Moreover, regardless of whether this system is shown to be a genuine massive cluster, the concentration of EROs is likely a sign that this field is an uncommonly overdense region at high redshift, which warrants follow-up studies.  Further observations are needed to develop a full physical picture of this field.  A much larger sample of redshifts, inferred from multicolor photometric redshifts and directly measured from deep spectroscopy, will reveal the number of luminous galaxies at this redshift.  An expanded spectroscopic sample will also be valuable to measure the galaxy velocity distribution and to search for dynamical sub-structure.  X-ray observations will allow us to look for hot intracluster gas, a sign of a deep gravitational potential, and also for mass sub-structure. This issue can also be addressed with radio observations to search for a Sunyaev-Zeldovich decrement.  Extending galaxy cluster studies to $z>1$ has special importance since current hierarchical formation theories predict massive galaxies assemble from smaller sub-units during this epoch \\citep[e.g.,][]{1998MNRAS.294..705K}, so we can potentially test these models by witnessing the formation process in situ.  Measuring the high-redshift cluster abundance using wide-area surveys can provide strong tests of cosmological models \\citep{1996MNRAS.282..263E,1997ApJ...485L..53B}.  However, even before a large number of such clusters have been found, the few known to date are valuable sites to investigate the processes which drive the evolution of the oldest stellar populations, galaxies, and galaxy clusters. Furthermore, these clusters can also be used as testing grounds for $z>1$ cluster-search strategies, which by necessity are derived from extrapolating the known properties of lower-redshift clusters.  It may be that we have to abandon some common precepts formed from studying local rich clusters in order to develop a complete understanding of high-redshift clusters and their constituent galaxies at an epoch of less than half the current age of the Universe."
    },
    "0002/astro-ph0002396_arXiv.txt": {
        "abstract": "{ We report spectral time series of the late O-type runaway supergiant HD 188209. Radial velocity variations of photospheric absorption lines with a possible quasi-period $\\sim$ 6.4 days have been detected in high-resolution echelle spectra. Night-to-night variations in the position and strength of the central emission reversal of the \\Ha~profile occuring over ill-defined time-scales have been observed. The fundamental parameters of the star have been derived using state-of-the-art plane--parallel and unified non-LTE model atmospheres, these last including the mass-loss rate. The derived helium abundance is moderately enhanced with respect to solar, and the stellar masses are lower than those predicted by the evolutionary models. The binary nature of this star is not suggested either from {\\it Hipparcos\\/} photometry or from radial velocity curves. } ",
        "introduction": "Runaway O stars have been defined as a group by Blaauw (1961), who introduced the term {\\it runaway} to describe the space motions of AE Aur and $\\mu$ Col. Blaauw (1961) has also suggested that such stars were ejected in the breakup of binary systems in supernova explosions by their companions. In later evolutionary stages, the initial secondary appears as a most massive star and transfers matter to the compact companion (the initial primary) making the system appear as a massive X-ray binary (van den Heuvel 1976). Given the possibility of the binary nature of runaway stars, it appears to be an important task to measure the radial velocity (RV) variations of the photospheric lines. Systematic searches for RV variations have been made in order to assess the binary frequency of O stars (e.g. Garmany, Conti \\& Massey 1980; Stone 1982, Gies 1987). In many cases the amplitude of RV variations is quite large, and the additional presence of a clear periodicity immediately suggests a binary nature for the system. However, there are stars which show more complicated RV curves, and the interpretation of their spectral variability is not straightforward. HD 188209 (O9.5Iab) is one of those objects. Garmany et al. (1980) have concluded from three spectra that this star is probably not a binary, and that the RV variations must be attributed to atmospheric motions. This conclusion was supported by Musaev \\& Chentsov (1988). However, based on 21 measurements Stone (1982) has concluded that HD 188209 can be considered as a spectroscopic  binary with  a period 57 days and small semiamplitude. More recently, Fullerton, Gies \\& Bolton (1996) included HD 188209 in a large sample of stars investigated on the presence of line profile variability (LPV) and found LPVs only in \\HeI~5876 \\AA. However, they did not flag HD 188209 as a velocity variable (their Table 10).  The binarity of many O supergiants has been proposed recently by Thaller (1997). The fact that binaries have a higher incidence and an H$\\alpha$ emission strength in post-MS stages may indicate that wind interactions are a common source of emission in massive stars. In other words, even in cases where RV measurements are not available, the presence of \\Ha~emission in a spectrum could be linked with colliding winds. One needs to study orbital phase variations in the \\Ha~profile in order to be sure that the latter is due to colliding winds instead of some other mechanism. Note that HD 188209 is an X-ray source detected by {\\it ROSAT} (Berghoefer, Schmitt \\& Cassinelli 1996).  In this paper we focus on the high-resolution spectroscopic data of HD 188209. Our observations can possibly account for the small semi-amplitudes and eccentric orbits of this binary candidate since they have been accumulated at different periods over a long baseline. ",
        "conclusions": "Our target belongs to the group of stars for which the existence of a compact companion has been proposed in the literature. The task of disproving or confirming the  binary nature of the system can be tackled only if sufficiently accurate analysed observational data are available. In this paper we used state-of-the-art models of atmospheres to determine the fundamental parameters of HD 188209.  To establish the presence of a possible companion we have studied the RVs of absorption lines by combining them in different groups. Fourier analysis based on the iterative {\\sc clean} algorithm was used to search for periodic variability. Unfortunately the time coverage of runs 1--4 and 8--9 was too sparse to set constraints on their time-dependent behaviour. For this reason we first analysed a few runs separately and then utilized the {\\sc clean} algorithm to search for periods in a whole data set. The highest peak in the Fourier power spectrum was centered near the frequency 0.156 day$^{-1}$ (6.4 days).  The 6.4 days period can be due to the binary nature of the system if one assumes very small and unlikely values for the mass ratio (q$\\le$0.1). Taking the values derived in this article (\\M = 16.6, \\R =20.9) and assuming q=0.1, we obtain for the Roche radius  and for the major semi-axis of the binary orbit 15 \\Rsun~and 25 \\Rsun, respectively. This simple estimate shows that an O supergiant can hardly fit within the orbit because its Roche radius would be less than the stellar radius.  Even if it would fit, the tides in such a tight binary would be very strong making the star to speed up quickly until the rotation period matches the orbit. Even if we assume that the system is very young, it's hard to explain that the putative orbital period is 2 times shorter than the rotational period (13 days). The second difficulty with the binary interpretation comes from the variability of H$\\alpha$. A TVS (temporal variance spectrum, showing the extent and distribution of statistically significant profile variability) has been computed recently (Baade 1998b, Kaper et al. 1998) for 15 spectroscopic binaries and it was found that all they show a characteristic double-peaked profile. This is due to two H$\\alpha$ absorption/emission profiles moving in a composite spectra. In our case the H$\\alpha$ profile is splitted because of the central emission coming from the lower wind. The last argument comes from the clear relations between excitation energies and radial velocity amplitude and excitation energy and mean radial velocities and from the model atmosphere atmosphere calculations. The latter is a good discriminant between internal variations (pulsations \\& wind instabilities) and Keplerian motions. In a binary system one would expect all lines to have the same amplitude independent on their TEE.  It has been known for a long time (Abt 1957) that the quasi-periodicity in hot supergiants might be ascribed to radial pulsations. A simple relation (Burki 1978; de Jager 1980) can be used to estimate the period of radial pulsation, \\begin{equation} \\log P_{\\rm fund} = 10.93 - 0.5\\log (M/M_{\\odot}) - 0.38M_{\\rm bol} - 3 \\log T_{\\rm eff} \\end{equation} Using the values of parameters obtained in Section 4 we arrive at $P_{\\rm fund}$=1.75 d. Note that the form of the relation (1) depends on the stellar evolutionary models and the input parameters; both are subject to large errors. In particular, note that we found no large differences in the parameters determined with plane--parallel and unified model atmospheres. Nevertheless, Levy et al. (1984) have pointed out that periods a factor of 1.5 longer than the corresponding periods of the radial pulsations can be ascribed to non-radial pulsations. This means that a factor of two difference between the evolutionary and the spectroscopic masses can easily result in the mis-identification of the pulsating mode. Another difficulty has been pointed by the referee of the article. A more sophisticated approach shows (Unno et al. 1979) that f-mode pulsation (which is the lowest-frequency mode supported by radial pulsation) periods are about 10 times larger than the one suggested by a period-luminosity relation.  In any case, the theoretical period of 1.75 days is very close to the Nyquist frequency of our data which means that we have a little chance to identify it in our data set even if it exists.  Our data not allow to distinguish between pulsations and stochastic variations of the stellar wind. It is also quite possible that we have a combination of both effects.  Note that the projected rotational period of this star ($\\sim$ 13 d) is much longer than any of the quasi-periods found in this paper (but of course our runs do not cover a whole rotation cycle). The surface features (if any) will always be visible on the projected disc of the star independently of the inclination angle. Thus, any periods due to the rotation of these features must correspond directly to the rotation period. We do not find any peaks in the power spectra at $\\sim$13 d and this leads us to discard rotational modulation as a possible explanation of the RV variations reported here.  The quality and the sampling of our data do not allow a careful study of the line asymmetries, moving components (if later exists) and/or long-term spectroscopic variability to be made. It is quite possible that the non-sinusoidal character of the RV curve for 6.4 days period (Fig 8) is caused by some disturbances due to the NRPs and/or moving features in the profiles plus any stochastic instabilities of the wind. New monitoring with much higher S/N may allow NRPs, multimode pulsations and clearly separate a sinusoidal curve of the radial pulsations to be revealed. However, we found convincing evidence that the atmospheric motions cannot be ascribed solely to Keplerian motions and probably are not of a binary origin."
    },
    "0002/astro-ph0002304_arXiv.txt": {
        "abstract": "The results obtained by means of the photometric approach to the study of $\\delta$ Sct stars are extensively discussed. The different frequency contents of the three best candidates for asteroseismological studies (FG Vir, 4 CVn and XX Pyx) are presented and compared; the importance of the amplitude variations and of the combination terms is emphasized. The analysis of other multiperiodic variables shows how a large variety of nonradial modes are excited; in some cases, modifications of the power spectrum can be observed over a few years and new modes can be seen to grow. Among monoperiodic pulsators, constant as well as variable amplitudes can be observed. The difficulty of identifying an oscillation yielding quantum numbers is emphasized; the possibilities offered by $\\delta$ Sct stars belonging to binary systems and open clusters are discussed. In this respect, combining the photometric and the spectroscopic approaches could lead to a solution. A comparison is also made between low- and high-amplitude pulsators, finding similarities.  The use of a reliable Period--Luminosity--Colour relationships toward the shortest periods can greatly help mode identification in the galactic stars; moreover, it could provide an independent verification of extragalactic distances. ",
        "introduction": "$\\delta$ Sct variables are now a well-defined class of stars. They are located on or just above the zero-age main sequence, in the lowest part of the classical instability strip. $\\delta$ Sct stars have masses between 1.5 and 2.5 M$_{\\sun}$ and they are close to the end of the core hydrogen burning phase (Breger \\& Pamyatnykh 1998). The presence of convective zones and related phenomena such as convective overshooting, make them very interesting objects for the understanding of stellar evolution. The investigation of  pulsational properties of pre-main sequence stars allowed their instability strip in the H-R diagram to be defined (Marconi \\& Palla 1998); some of these stars showing $\\delta$ Sct variability were discovered (Kurtz \\& Muller 1999).  Photometric monitoring is the most practiced approach to study the properties of $\\delta$ Sct stars and several stars have been deeply investigated. However, the spectroscopic approach tells us that many modes that are not photometrically detectable are actually excited. Rotation acts as an important factor in the increase of the number of excited modes.  The observed frequencies are between 5 and 35~cd$^{-1}$ and multiperiodicity is very common; only a few stars show a monoperiodic behaviour above the current limit of the detectable amplitude from ground, i.e., $\\sim$1~mmag. The observed modes are in the domain of pressure ($p$) modes; there is no observational evidence that gravity ($g$) modes are excited in $\\delta$ Sct stars, even if some cases are suggested. At the moment, $g$-modes seem to be present only in the $\\gamma$ Dor stars, which in turn do not show $p$-modes.  The great observational effort made by several teams allows us to handle a well-defined phenomenological scenario of the $\\delta$ Sct variability. This contribution tries to summarize the results obtained in the past years by means of extended photometric time series. The paper is  structured as follows: \\begin{list} { } { } \\item 2.~The best candidates for asteroseismological studies \\begin{list} { } { } \\item 2.1~FG Vir: 24 independent modes \\item 2.2~XX Pyx: strong and rapid amplitude changes \\item 2.3~4 CVn: presence of combination terms and amplitude variations \\item 2.4~Comparison between FG Vir, XX Pyx and 4 CVn \\end{list} \\item 3.~Other stars studied by the Merate Group \\begin{list} { } { } \\item 3.1~44 Tau: variable amplitude and recurrent ratio 0.77 \\item 3.2~BH Psc: variable amplitude and rich pulsational content \\item 3.3~V663 Cas: growth of new modes \\item 3.4~The help of the spectroscopic approach \\end{list} \\item 4.~Monoperiodic pulsators \\item 5.~$\\delta$ Sct stars in binary systems \\item 6.~$\\delta$ Sct stars in open clusters \\item 7.~The frequency content of high amplitude $\\delta$ Sct stars \\item 8.~Summing-up and Conclusions \\item 9.~The future: the exportation of the results on galactic stars to \\\\ \\hspace*{4truemm} extragalactic research \\end{list} In what manner the photometric results can be used to identify modes (i.e., to classify the oscillation in terms of quantum numbers $n$, $\\ell$ and $m$) is discussed  by Garrido (2000). Some improvements, both observational and theoretical, are probably necessary to make new, substantial steps forward. A better connection between theory and observation will allow us to really progress in the asteroseismology of these stars. ",
        "conclusions": "In our travel through the observational scenario we met many objects which teach us something about the pulsation of $\\delta$ Sct stars. The high frequencies shown by XX Pyx, the intermediate frequencies displayed by FG Vir and the low frequencies found in the case of 4 CVn  are leading toward the idea that the mode excitation is really complex and unpredictable. In this respect, the close similarity between 4 CVn and HD~2724 is comfortable. Probably the multiperiodic behaviour of the high-amplitude $\\delta$ Sct star V974 Oph can also be seen as a simplification of the phenomenology, demonstrating that we can observe a multimode excitation even when  a large pulsational energy is involved. To complete the similarities in the opposite direction, we observe monoperiodic pulsators with a very small amplitude (HD 19279 and $\\beta$ Cas). What the excited mode is should be investigated in order to understand if a  selection effect acts for these low-amplitude monoperiodic stars and what difference there is with a star such as AZ CMi, which displays an asymmetrical light curve.  The amplitude variations are observed in a large variety of stars, both multiperiodic (XX Pyx, 4 CVn, 44 Tau) and monoperiodic (28 And and BF Phe).  The observations of mode growth in the light curve of a relatively simple pulsator such as V663 Cas clarifies what can happen in much more complicated ones: some modes can be damped and then re-excited. There are some observational facts which make this explanation preferable even if the model of a beating between two close frequencies with similar, constant amplitude cannot be ruled out.  In this context the continuous survey of $\\delta$ Sct stars which will be performed by space missions could greatly improve the situation; in the case of intrinsic variations, the time-scale of such variations could be determined.  \\smallskip  The better focusing of the phenomenological scenario is not the only result.  The effort made by the observers in the last years allowed us to detect a large number of frequencies in stars such as FG Vir, XX Pyx, 4 CVn and BH Psc. By means of multisite campaigns we are able to detect terms with amplitudes less than 1 mmag; in these conditions ground-based observations can be successfully complementary to space missions. The mode identification techniques are based on both the phase shifts and amplitude ratios of the light and colour curves and the synergic approach performed by considering spectroscopic curves. Their full exploitation should guarantee an important role for our researche in stellar physics."
    },
    "0002/astro-ph0002418_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 present the Catalog of microlensing events detected toward the Galactic bulge in three observing seasons, 1997--1999, during the OGLE-II microlensing survey. The search for microlensing events was performed using a database of about $4\\cdot10^9$ photometric measurements of about 20.5 million stars from the Galactic bulge. The Catalog comprises 214 cases of microlensing events found in the fields covering about 11 square degrees on the sky and distributed in different parts of the Galactic bulge. The sample includes 20 binary microlensing events, 14 of them are caustic crossing. In one case a double star is likely lensed.  We present distribution of the basic parameters of microlensing events and show preliminary rate of microlensing in different regions of the Galactic bulge. The latter reveals clear dependence on the Galactic coordinates. The dependence on $l$ indicates that the majority of lenses toward the Galactic bulge are located in the Galactic bar. Models of the Galactic bar seem to reasonably predict the observed spatial distribution of microlensing events in the Galactic bulge.  All data presented in the Catalog and photometry of all events are available from the OGLE Internet archive. }{~} ",
        "introduction": "During the past couple of years microlensing proved to be  a new and potentially very powerful tool of modern astrophysics. Originally proposed by Paczy{\\'n}ski (1986, 1991) as a method of searching for dark matter in the Galaxy it has been used for such different applications as searching for planets, determination of parameters of stellar atmospheres, studies of Galactic structure and many others.  After the original reports on discovery of the first cases of microlensing in September 1993 (toward the LMC: MACHO survey -- Alcock \\etal 1993, EROS survey -- Aubourg \\etal 1993; toward the Galactic bulge:  OGLE survey -- Udalski \\etal 1993) much observing work was done to convince the astronomical community on the potentials of the newly discovered class of events. Soon more cases of classical microlensing in the Galactic bulge were announced (Udalski \\etal 1994a, Alcock \\etal 1995a), first cases of \"exotic microlensing\" like binary microlensing (Udalski \\etal 1994b) or events with parallax effect (Alcock \\etal 1995b) were found. Also first estimates of the observed optical depth to the Galactic bulge were published (Udalski \\etal 1994c, Alcock \\etal 1995a, 1997a) indicating that it is much larger than that predicted from modeling. Many theoretical interpretations of these intriguing results followed (see review by Paczy{\\'n}ski 1996). First interpretation of results of observations toward the LMC was also published (Alcock \\etal 1996a).  Another important step was development and implementation of the so called alert systems (the Early Warning System, EWS, of the OGLE survey -- Udalski \\etal 1994d and MACHO Alert system -- Alcock \\etal 1996b) which allow to detect the microlensing phenomena when an event is in progress. This step changed the observing strategy of microlensing surveys. The ability of detection in real time  made it possible to perform follow-up observations, both photometric and spectroscopic, of many events. New kind of microlensing studies, \"follow-up\" projects concentrated on high time resolution observations of events discovered by survey projects, were formed (\\eg PLANET -- Albrow \\etal 1998; MPS -- Rhie \\etal 1999). From 1995 on, the vast majority of microlensing events have been detected by the alert systems.\\footnote{\\noindent Information on microlensing events in progress can be found:\\newline OGLE-EWS: {\\it http://www.astrouw.edu.pl/\\~{}ogle} \\newline MACHO (1995--1999): {\\it http://darkstar.astro.washington.edu/} \\newline EROS: {\\it http://www-dapnia.cea.fr/Spp/Experiences/EROS/alertes.html} \\newline PLANET: {\\it http://www.astro.rug.nl/\\~{}planet/} \\newline MPS: {\\it http://bustard.phys.nd.edu/MPS/} }  The microlensing field of astrophysics matured rapidly and entered the second phase of extensive observations to largely increase statistic of collected microlensing events. Both EROS and OGLE surveys considerably increased their observing capabilities in 1996. First microlensing events were discovered in other lines of sight -- toward the SMC (MACHO -- Alcock \\etal 1997b) and in the Galactic disk (EROS -- Derue \\etal 1999) even as far as $70^{\\circ}$ from the Galactic center (OGLE -- Mao 1999). Many interesting results were published by follow-up teams (PLANET -- Albrow \\etal 2000a,b, MPS -- Bennett \\etal 1999a). Up to now one can estimate the total number of registered microlensing events to be about 500.  In this paper we present the catalog of microlensing event candidates detected during the second phase of the Optical Gravitational Lensing Experiment (OGLE-II) in three observing seasons 1997--1999. Although during the OGLE-II phase observations are conducted in many lines of sight only two events were discovered in the directions other than the Galactic bulge so far.\\footnote{1999-CAR-01 and 1999-LMC-01 see {\\it http://www.astrouw.edu.pl/\\~{}ogle} for more information on these events.} Therefore we decided to limit our catalog to the Galactic bulge events only. Presented microlensing events were extracted from the OGLE-II photometric databases with a technique similar to that applied in the EWS alert system. Great care was paid to achieve the highest possible completeness of the catalog. The Catalog comprises 214 cases of microlensing toward the Galactic bulge.  The main goal of the paper is to provide the astronomical microlensing community with the large set of photometric data of microlensing events for further analysis. Photometry of all objects is available from the OGLE Internet archive (see Section~5). Beside of the typical single mass microlensing cases the Catalog includes several cases of \"exotic microlensing\", like binary microlensing events etc.  The sample of presented events is already large enough that in spite of possible incompleteness the distribution  of microlensing parameters can be studied. In particular we present the first preliminary distribution of the rate of microlensing in different parts of the Galactic bulge. Studies of such a distribution can provide important constraints on the origin of the Galactic bulge microlensing (bar \\vs disk) and in general on the structure of the Galaxy. ",
        "conclusions": "Two hundred fourteen microlensing event candidates were detected in the OGLE-II fields during the observing seasons 1997--1999. Significant number of events from 1998--1999 seasons, \\ie when the EWS alert system was implemented, were detected in real time. Many of these events were then followed up by other groups for detailed study (\\eg 1998-BUL-14, Albrow \\etal 2000a).  Because part of the OGLE-II fields overlap with those observed by the MACHO team several events were discovered by both teams. Cross-identification of those discovered by alert systems can be found on the WWW alert pages of both teams while for the remaining objects (in particular all 1997 season events) it can be done using accurate equatorial coordinates provided in this paper and on the MACHO WWW alert page.  The sample of OGLE-II microlensing events contains several cases of evident lensing caused by binary object. The sample of characteristic caustic crossing events consists of 14 events. In five additional cases it is possible that the observed light variations are also caused by binary lens. In two cases, BUL$\\_$SC39 259656 and BUL$\\_$SC40 434222, two separate microlensing like episodes were detected. These cases might be explained either as due to  binary microlensing without caustic crossing or double lensed star. In the latter case the time scales of both episodes should be equal what can discriminate between those two possibilities. Time scales of BUL$\\_$SC39 259656 episodes are equal to about 15 and 27~days suggesting that this case is a binary microlensing. The second brightening episode of BUL$\\_$SC40 434222 started at the end of 1999 season but observations collected up to the moment of writing this paper confirm that it can be of microlensing origin (see OGLE EWS WWW page). In this case the time scales are within errors the same ($t_0\\approx 130$~days) suggesting double source star. However, the second episode is still in the rising part of its light curve.  \\begin{figure}[htb] \\psfig{figure=fig2.ps,bbllx=60pt,bblly=50pt,bburx=505pt,bbury=405pt,width=12cm,clip=} \\FigCap{Distribution of the {\\it I}-band baseline magnitude of lensed stars.} \\end{figure}  The rate of the binary caustic crossing microlensing toward the Galactic bulge observed in our sample of all microlensing events is equal to $\\approx 6.5$\\% while that of all binary microlensing $\\approx 9.3$\\%. It is worth noting that this is much more than reported by the MACHO team (Alcock \\etal 1999). On the other hand Mao and Paczy{\\'n}ski (1991) in their classical paper on binary microlensing predicted about 10\\% rate of strong binary microlensing events. The agreement between the predicted and observed rates is very good.  Among many interesting cases of microlensing presented in the Catalog we draw attention to the object, BUL$\\_$SC5 244353. Unfortunately, only falling part of its light curve was covered during the entire observation period. While the shape of the light curve strongly resembles that of microlensing event, it cannot be excluded that the object is a variable star. If, however, its brightening was caused by microlensing, the time-scale of the event was very long. The event would closely resemble those reported by Bennett \\etal (1999b) which are supposed to be caused by black hole lenses.   \\begin{figure}[htb] \\psfig{figure=fig3.ps,bbllx=60pt,bblly=50pt,bburx=505pt,bbury=405pt,width=12cm,clip=} \\FigCap{Magnification at maximum, $A_{\\rm max}$, and minimum impact parameter, $u_{\\rm min}$, as a function of the {\\it I}-band baseline magnitude. } \\end{figure}  The statistic of collected events toward the Galactic bulge is becoming significantly larger after each observing season. While an accurate analysis requires additional information like  efficiency of detection, the number of presented microlensing events is already so large  that interesting correlations can be shown.  Efficiency of detection of microlensing events depends at least on two factors: time sampling of the microlensing light curve and brightness of the lensed star. Typical sampling of one/two observations per night limits detection to events with sufficiently long time-scale. For fainter lensed stars the $3\\sigma$ detection threshold is larger, therefore only larger magnification events can be triggered. In Fig.~2 we plot distribution of the {\\it I}-band baseline magnitude of lensed stars from our sample. The bins are 0.3~mag wide. Dotted line presents the distribution of all events while the bold solid line the one of subsample of events with magnification $A_{\\rm max}>1.5$. The shape of both distributions indicate that the entire sample of microlensing events is reasonably complete up to  $I\\approx18.0$~mag while subsample of events with $A_{\\rm max}>1.5$ up to $I\\approx18.8$~mag.  Selection effect due to brightness of the lensed star can also be assessed from Fig.~3 which shows magnification at the maximum,$A_{\\rm max}$, and minimum impact parameter, $u_{\\rm min}$, plotted against the {\\it I}-band baseline magnitude. It is clearly seen from Fig.~3 that our sample is quite complete for stars of $I<18$~mag and $A_{\\rm max}>1.3$, \\ie $u<1.0$. As can be expected for fainter stars only higher magnification events were triggered. Nevertheless at $I=19$~mag the limit of reasonable completeness is still $A_{\\rm max}\\approx1.5$ ($u_{\\rm min}\\approx0.8$).  \\begin{figure}[htb] \\psfig{figure=fig4.ps,bbllx=60pt,bblly=50pt,bburx=505pt,bbury=405pt,width=12.5cm,clip=} \\FigCap{Distribution of the maximum magnification $A_{\\rm max}$.} \\end{figure}  Fig.~4 presents distribution of maximum magnification of our entire sample. This parameter ranges from as low as 1.1 up to about 50 (in a few cases the single point mass model predicts extremely large magnifications but because of lack of observations at the very maximum these values are not reliable.  In these cases we provide lower limit of magnification resulting from the brightest observation). Below $A_{\\rm max}=1.3$ incompleteness of our sample is larger -- only for brighter stars so low magnification events could be detected. However, the completeness becomes much higher for events with $A_{\\rm max}>1.5$. The number of events is gradually falling from $A_{\\rm max}\\approx1.7$ to $A_{\\rm max}=7$ with a long tail of single events with higher $A_{\\rm max}$.  \\begin{figure}[htb] \\psfig{figure=fig5.ps,bbllx=60pt,bblly=50pt,bburx=505pt,bbury=405pt,width=12.5cm,clip=} \\FigCap{Distribution of the Einstein radius crossing time $t_0$.} \\end{figure}  Fig.~5 shows the distribution of the Einstein radius crossing time, $t_0$, for our sample of microlensing events. This parameter is also likely to be affected by incompleteness resulting from the sampling of the light curve. Although we have not performed yet detailed analysis of the dependence of detection efficiency on the event time-scale  for our entire data set, preliminary tests performed on 1997 databases of constant stars in similar manner as in Udalski \\etal (1994c) indicate that the region of reasonable efficiency is extended toward shorter time scales  as compared to the OGLE-I phase. For events with $t_0>8$ days efficiency of detection becomes relatively flat, so we may expect that the distribution of $t_0$ is also relatively complete for events longer than that limit. The distribution of $t_0$ peaks at $t_0\\approx17$~days with a long tail of longer time-scale events. One should also note  a small excess of events with $t_0\\approx 50$~days. It was marginally seen in the MACHO data (Alcock \\etal 1997a). The time-scale is plotted as a function of the {\\it I}-band baseline magnitude in Fig.~6. No evident correlation is seen in this plot.  \\begin{figure}[htb] \\psfig{figure=fig6.ps,bbllx=60pt,bblly=50pt,bburx=505pt,bbury=405pt,width=12.5cm,clip=} \\FigCap{The Einstein radius crossing time, $t_0$, as a function of the {\\it I}-band baseline magnitude. } \\end{figure}   \\begin{figure}[htb] \\psfig{figure=fig7.ps,bbllx=35pt,bblly=50pt,bburx=505pt,bbury=405pt,width=12.5cm,clip=} \\FigCap{Spatial distribution in the Galactic bulge of the short (open circles), medium (open triangles) and long (filled squares) time-scale microlensing events.} \\end{figure}  The spatial distribution of events with different time-scales is presented in Fig.~7. The total sample was divided into three sub-samples of short ($t_0<20$~days), medium ($20<t_0<40$~days) and long time-scale ($t_0>40$~days) events. Location on the sky of these three samples in the Galactic coordinates ($l,b$) is plotted with different symbols in Fig.~7. If the different time-scale events were to be caused by different populations of lensing objects one could expect some differences in the distribution of our three sub-samples on the sky. However, this does not seem to be true -- the distribution of all our sub-samples is rather similar. One should be, however, aware of still small statistic of events in the regions at larger $|l|$.  \\renewcommand{\\arraystretch}{1.0} \\renewcommand{\\TableFont}{\\scriptsize} \\MakeTableSep{lrcrr}{12.5cm}{Number of microlensing events in the OGLE-II fields} { \\hline \\noalign{\\vskip3pt} \\multicolumn{1}{c}{Field} & $N_{\\rm OBS}^{{\\mu}{\\rm LENS}}$  & $N_{I<19.5}^{\\rm STARS}$ &\\multicolumn{1}{c}{$N^{{\\mu}{\\rm LENS}}$ per } \\\\ & & &\\multicolumn{1}{c}{$10^6$ stars}\\\\ \\noalign{\\vskip3pt} \\hline \\noalign{\\vskip3pt} BUL$\\_$SC1   &   1~~~~~ & 579010 &  1.7~~~~~\\\\ BUL$\\_$SC2   &   4~~~~~ & 650973 &  6.1~~~~~\\\\ BUL$\\_$SC3   &  12~~~~~ & 644230 & 18.6~~~~~\\\\ BUL$\\_$SC4   &  14~~~~~ & 644012 & 21.7~~~~~\\\\ BUL$\\_$SC5   &   5~~~~~ & 280867 & 17.8~~~~~\\\\ BUL$\\_$SC6   &   1~~~~~ & 289813 &  3.5~~~~~\\\\ BUL$\\_$SC7   &   1~~~~~ & 280966 &  3.6~~~~~\\\\ BUL$\\_$SC8   &   0~~~~~ & 205248 &  0.0~~~~~\\\\ BUL$\\_$SC9   &   0~~~~~ & 212660 &  0.0~~~~~\\\\ BUL$\\_$SC10  &   3~~~~~ & 222382 & 13.5~~~~~\\\\ BUL$\\_$SC11  &   0~~~~~ & 199964 &  0.0~~~~~\\\\ BUL$\\_$SC12  &   0~~~~~ & 312079 &  0.0~~~~~\\\\ BUL$\\_$SC13  &   1~~~~~ & 335321 &  3.0~~~~~\\\\ BUL$\\_$SC14  &   7~~~~~ & 414683 & 16.9~~~~~\\\\ BUL$\\_$SC15  &   7~~~~~ & 361358 & 19.4~~~~~\\\\ BUL$\\_$SC16  &   2~~~~~ & 435850 &  4.6~~~~~\\\\ BUL$\\_$SC17  &   2~~~~~ & 454107 &  4.4~~~~~\\\\ BUL$\\_$SC18  &   7~~~~~ & 556327 & 12.6~~~~~\\\\ BUL$\\_$SC19  &   3~~~~~ & 508643 &  5.9~~~~~\\\\ BUL$\\_$SC20  &   6~~~~~ & 692281 &  8.7~~~~~\\\\ BUL$\\_$SC21  &   4~~~~~ & 704858 &  5.7~~~~~\\\\ BUL$\\_$SC22  &   7~~~~~ & 503527 & 13.9~~~~~\\\\ BUL$\\_$SC23  &   6~~~~~ & 455213 & 13.2~~~~~\\\\ BUL$\\_$SC24  &   7~~~~~ & 410873 & 17.0~~~~~\\\\ BUL$\\_$SC25  &   1~~~~~ & 437199 &  2.3~~~~~\\\\ BUL$\\_$SC26  &   5~~~~~ & 446205 & 11.2~~~~~\\\\ BUL$\\_$SC27  &   6~~~~~ & 434129 & 13.8~~~~~\\\\ BUL$\\_$SC28  &   0~~~~~ & 258811 &  0.0~~~~~\\\\ BUL$\\_$SC29  &   2~~~~~ & 257335 &  7.8~~~~~\\\\ BUL$\\_$SC30  &  12~~~~~ & 606406 & 19.8~~~~~\\\\ BUL$\\_$SC31  &   7~~~~~ & 651611 & 10.7~~~~~\\\\ BUL$\\_$SC32  &   3~~~~~ & 680714 &  4.4~~~~~\\\\ BUL$\\_$SC33  &   4~~~~~ & 563023 &  7.1~~~~~\\\\ BUL$\\_$SC34  &   7~~~~~ & 743418 &  9.4~~~~~\\\\ BUL$\\_$SC35  &   5~~~~~ & 614900 &  8.1~~~~~\\\\ BUL$\\_$SC36  &   2~~~~~ & 659023 &  3.0~~~~~\\\\ BUL$\\_$SC37  &   8~~~~~ & 509135 & 15.7~~~~~\\\\ BUL$\\_$SC38  &   5~~~~~ & 572045 &  8.7~~~~~\\\\ BUL$\\_$SC39  &  15~~~~~ & 627658 & 23.9~~~~~\\\\ BUL$\\_$SC40  &   6~~~~~ & 366672 & 16.4~~~~~\\\\ BUL$\\_$SC41  &   8~~~~~ & 394514 & 20.3~~~~~\\\\ BUL$\\_$SC42  &   4~~~~~ & 436339 &  9.2~~~~~\\\\ BUL$\\_$SC43  &   4~~~~~ & 295349 & 13.5~~~~~\\\\ BUL$\\_$SC44  &  10~~~~~ & 183296 & 54.6~~~~~\\\\ BUL$\\_$SC47  &   0~~~~~ & 155703 &  0.0~~~~~\\\\ BUL$\\_$SC48  &   0~~~~~ & 156226 &  0.0~~~~~\\\\ BUL$\\_$SC49  &   0~~~~~ & 148155 &  0.0~~~~~\\\\ \\hline}  Table~3 lists the average number of stars searched for microlensing events in each of the Galactic bulge fields. In the second column the number of detected events in each field is provided. Because the number of searched stars is different by a factor of more than four in our fields, we normalized the observed number of microlensing events in each field to one million stars. Normalized number of events is listed in the last column of Table~3.    \\MakeTable{lrrr}{12.5cm}{Average number of microlensing events in the Galactic bulge} { \\hline \\noalign{\\vskip3pt} \\multicolumn{1}{c}{Line of sight} & \\multicolumn{1}{c}{$l$} & \\multicolumn{1}{c}{$b$} &\\multicolumn{1}{c}{$N$ per}\\\\ \\multicolumn{1}{c}{BUL$\\_$SC} & & & $10^6$ stars \\\\ \\hline \\noalign{\\vskip3pt} 5+44\t&\t$ -0.33$\t& $-1.26$ &\t36~~~~~\\\\ 3+37\t&\t$  0.05$\t& $-1.83$ &\t17~~~~~\\\\ 4+39\t&\t$  0.48$\t& $-2.11$ &\t23~~~~~\\\\ 22+23\t&\t$ -0.38$\t& $-3.15$ &\t14~~~~~\\\\ 6+7\t&\t$ -0.20$\t& $-5.80$ &\t 4~~~~~\\\\ 40+41\t&\t$ -2.89$\t& $-3.20$ &\t18~~~~~\\\\ 24+25\t&\t$ -2.38$\t& $-3.46$ &\t10~~~~~\\\\ 26+27\t&\t$ -4.91$\t& $-3.51$ &\t13~~~~~\\\\ 28+29\t&\t$ -6.70$\t& $-4.52$ &\t 4~~~~~\\\\ 47+48+49&\t$ -11.21$\t& $-2.88$ &\t 0~~~~~\\\\ 1+38\t&\t$  1.02$\t& $-3.52$ &\t 5~~~~~\\\\ 20+34\t&\t$  1.52$\t& $-2.43$ &\t 9~~~~~\\\\ 21+30\t&\t$  1.87$\t& $-2.75$ &\t13~~~~~\\\\ 31+32\t&\t$  2.28$\t& $-3.04$ &\t 8~~~~~\\\\ 35+36\t&\t$  3.10$\t& $-3.10$ &\t 6~~~~~\\\\ 2+33\t&\t$  2.29$\t& $-3.56$ &\t 7~~~~~\\\\ 18+19\t&\t$  4.03$\t& $-3.24$ &\t 9~~~~~\\\\ 42\t&\t$  4.48$\t& $-3.38$ &\t 9~~~~~\\\\ 16+17\t&\t$  5.19$\t& $-3.37$ &\t 4~~~~~\\\\ 12+13\t&\t$  7.85$\t& $-3.48$ &\t 2~~~~~\\\\ 10+11\t&\t$  9.69$\t& $-3.54$ &\t 7~~~~~\\\\ 8+9\t&\t$ 10.53$\t& $-3.88$ &\t 0~~~~~\\\\ 14+15\t&\t$  5.30$\t& $ 2.72$ &\t18~~~~~\\\\ 43\t&\t$  0.37$\t& $ 2.95$ &\t14~~~~~\\\\ \\hline}  Finally, because we typically observe two slightly overlapping driftscan fields in a given part of the Galactic bulge (see Fig.~1), we averaged the normalized numbers of events in such adjacent fields to increase statistic. Table~4 lists our 24 lines of sight in the Galactic bulge with their Galactic coordinates and the average number of observed microlensing events per one million stars during seasons 1997--1999 (the total span of presented observations is about 940 days, \\ie 2.58~years).  \\begin{figure}[htb] \\psfig{figure=fig8.ps,bbllx=15pt,bblly=150pt,bburx=565pt,bbury=430pt,width=12.5cm,clip=} \\FigCap{Rate of microlensing events in the Galactic bulge. The numbers in the Galactic coordinates grid correspond to the number of microlensing events per one million stars observed during three bulge seasons 1997--1999 (2.58~years).} \\end{figure}  Fig.~8 presents the number of events per one million stars observed in 24 lines of sight in the Galactic bulge in years 1997--1999. The number at given ($l$,$b$) indicates the normalized number of events observed in a given direction. Of course, one should be aware that the presented numbers are not the true optical depth. Accurate estimate of the optical depth would require precise determination of detection efficiency, assessment  of the blending effect and contribution of the Galactic disk stars to the stars searched for microlensing events. On the other hand one can expect that these factors are in the first approximation similar in so uniformly observed fields. Also the spatial distribution of the time-scale of events is similar. Therefore the normalized number of events is likely a crude approximation of the optical depth.  As one can expect the number of events is a strong function of the Galactic latitude. It falls by almost an order of magnitude when $b$ changes from $b=-1\\zdot\\arcd3$ to $b=-6\\arcd$ (at $l\\approx0\\arcd$). It also changes with the Galactic longitude. The numbers of events observed in the fields at $b\\approx-3\\zdot\\arcd5$ indicate that there is a  clear dependence on $l$ -- the number of events falls by a factor of 2--4 at $|l|\\approx 10\\arcd$ as compared to $l=0\\arcd$. There is a noticable asymmetry with larger number of events at negative $l$. The clear dependence of the number of microlensing events on $l$ strongly suggests that the majority of microlensing events are caused by lenses located in the Galactic bar, inclined to the line of sight toward the Galactic center,  rather than in the Galactic disk.  At positive $b$  the number of microlensing events in the line of sight located  at $l=0\\zdot\\arcd4, b=3\\arcd$ is  consistent with that observed at negative $b$ but one can notice possible excess of events at $l=5\\zdot\\arcd3, b=2\\zdot\\arcd7$.  Looking at the numbers presented in Fig.~8 and Table~4 one should be aware of the possible bias resulting from somewhat different time sampling of some fields. While the numbers of observations in most fields are similar there are six fields with about 40\\% larger number of analyzed frames (see Table~1). These were the fields located closest to the Galactic center and they were observed with frequency of 3--5 observations per night during a part of the 1997 season for ultra-short time events. No such events (time scale $t_0<2$~days) were, however, found. Nevertheless, the number of detected microlensing events can be larger in more  frequently observed fields resulting in  somewhat overestimated  ratios between the number of microlensing events close to the Galactic center and in other directions.  We estimated the possible magnitude of this effect by comparison of the median time scale, $t_0$, of events from the least-, medium- and most-frequently observed fields. One could expect that for the most frequently sampled fields, the efficiency of detection is much higher for short time scale events and the median $t_0$ of that subsample should be significantly shorter than that of the remaining subsamples if the bias is strong. However, the median $t_0$ is equal to 17.0, 21.5 and 19.7~days for the most-, medium- and least-frequently sampled fields, respectively. The differences are small indicating that the bias is also small and can  be in the first approximation neglected. It would be certainly much stronger if much larger part of events had very short time scales.  \\begin{figure}[htb] \\psfig{figure=fig9.ps,bbllx=60pt,bblly=50pt,bburx=505pt,bbury=405pt,width=12.5cm,clip=} \\FigCap{Correlation between the observed number of microlensing events and the model optical depth.} \\end{figure}  It is interesting to compare our empirical results with the modeling predictions. We used the map of the optical depth for the Galactic bar calculated by Stanek \\etal (1997) to derive the optical depth in our lines of sight. Contribution of the Galactic disk in these directions was taken from Kiraga (1994) and added to the contribution of the Galactic bar. Fig.~9 shows the relation between the model optical depth and the observed number of microlensing events per one million stars in our lines of sight. The correlation between both values is clearly seen. The slope of the relation is equal to $0.86\\pm0.11$ giving a crude calibration between the observed numbers and the optical depth. However, one should also note the potential problem -- the linear relation does not cross the (0,0) point. This may indicate that the ordinate in Fig.~9 is not that simply related with the optical depth and the observed normalized number of events is a function not only of  the optical depth but also of an additional factor like, for instance, spatial dependent efficiency of detection.  Another possibility is that the model underestimate significantly the optical depth close to the Galactic center. In general, however, Figs.~8 and 9 indicate that models of the bar might provide a reasonable approximation of the Galactic bulge structure (see also for example Evans 1994, Zhao and Mao 1996, Grenacher \\etal 1999)  and that the microlensing will be a very powerful tool in further constraining the Galactic bulge properties when significantly larger statistic of events is collected.  While Fig.~8 can provide a first outlook on the distribution of microlensing events in the Galactic bulge one should be aware that this is only the first approximation. The statistic of microlensing events, in particular in the fields with larger $|l|$, is still small. Observations with the same set-up will be continued during the next observing season (2000) providing about 60--70 new cases of microlensing events. After that a large instrumental upgrade to the OGLE-III phase is planned by implementation of a new mosaic CCD camera. This will increase the number of discovered events by a factor of 3--5 leading to fast increase of the statistic of Galactic bulge events. Also larger area of the Galactic bulge will be monitored. It is also planned to reanalyze the OGLE-II photometric data with new techniques like image subtraction method. It would allow to get rid of blending effect uncertainties. By limiting to well defined population of  the Galactic bulge stars like for instance red clump giants it will be possible to provide much more accurate information on the optical depth distribution in the Galactic bulge in the near future.  The Catalog of Microlensing Events in the Galactic Bulge and all photometric data presented in this paper are available now to the astronomical community from the OGLE Internet archive:  \\begin{center} {\\it http://www.astrouw.edu.pl/\\~{}ogle} \\\\ {\\it ftp://sirius.astrouw.edu.pl/ogle/ogle2/microlensing/gb/}\\\\ \\end{center} or its US mirror \\begin{center} {\\it http://www.astro.princeton.edu/\\~{}ogle}\\\\ {\\it ftp://astro.princeton.edu/ogle/ogle2/microlensing/gb/}\\\\ \\end{center}  \\Acknow{We would like to thank Prof.\\ Bohdan Paczy\\'nski for many discussions and help at all stages of the OGLE project.  The paper was partly supported by  the Polish KBN grants 2P03D00814 to A.\\ Udalski, 2P03D00916 to M.\\ Szyma{\\'n}ski, and 2P03D00717 to K.\\ \\.Zebru\\'n. Partial support for the OGLE  project was provided with the NSF grant AST-9820314 to B.~Paczy\\'nski.}"
    },
    "0002/astro-ph0002132_arXiv.txt": {
        "abstract": "We present the results of broad-band B and I imaging observations for a sample of 88 Seyfert galaxies (29 Seyfert 1's and 59 Seyfert 2's), selected from a mostly isotropic property, the flux at 60$\\mu$m. We also present the B and I imaging results for an additional sample of 20 Seyfert galaxies (7 Seyfert 1's and 13 Seyfert 2's), selected from the literature and known to have extended radio emission.  The I band images are fitted with ellipses to determine the position angle and ellipticity of the host galaxy major axis. This information will be used in a future paper, combined with information from radio observations, to study the orientation of radio jets relative to the plane of their host galaxies (Kinney et al. 2000). Here we present surface brightness profiles and magnitudes in the B and I bands, as well as mean ellipticities and major axis position angles. ",
        "introduction": "\\label{sample}  \\subsection{60$\\mu$m sample}  In order to avoid selection effects as much as possible, we have chosen a sample from a mostly isotropic property, the flux at 60$\\mu$m. According to the torus models of Pier \\& Krolik (1992), which are the most anisotropic and hence the most conservative models, the circumnuclear torus radiates nearly isotropically at 60$\\mu$m.  Our sample includes 88 Seyfert galaxies (29 Seyfert 1's and 59 Seyfert 2's), which correspond to all galaxies from the de Grijp et al. (1987, 1992) sample of warm IRAS galaxies with redshift z$\\leq0.031$. The galaxies in this sample were selected based on the quality of the 60$\\mu$m flux, Galactic latitude $|b|>20^{\\circ}$, and 25$\\mu$m$-60\\mu$m color in the range $-1.5<\\alpha(25/60)<0$, chosen to exclude starburst galaxies as much as possible. The candidate AGN galaxies were all observed spectroscopically (de Grijp et al. 1992) to confirm their activity class as being Seyfert 1 or Seyfert 2 and {\\it not} a lower level of activity such as starburst or LINER. The distance limit of z$\\leq0.031$ is large enough to encompass a statistically significant number of objects yet close enough to ensure that radio features can be resolved.  Table 1 presents the galaxies in the de Grijp et al.  (1987) catalog, selected for our study. We list their catalog numbers, names, coordinates, the total exposure times in the B and I bands, and the observing runs in which the galaxies were observed.  \\subsection{Additional sample}  Parts of the study presented by Kinney et al. (2000) will also use an additional sample of 53 Seyfert galaxies selected from the literature. This sample comprises Seyferts known to have extended radio emission, used in previous studies (such as Schmitt et al. 1997; or Nagar et al. 1999) but which are not in the 60$\\mu$m sample. For 20 of these galaxies (7 Seyfert 1's and 13 Seyfert 2's) we were able to obtain B and/or I images during our observing runs. Table 2 gives the names of the galaxies, their coordinates, total exposure times in B and I bands and the observing run in which they were observed.  Some of the galaxies in the additional sample were used in previous papers, but we now consider that they should not be included in this analysis. The reasons to exclude them are the fact that they are in interacting systems, mergers, or the radio emission is not extended enough to allow a reliable measurement of the position angle of the jet. For these galaxies,  Column 8 (Comments) of Table 2 gives the reasons why they are excluded. ",
        "conclusions": "We presented B and I band images for a sample of 88 Seyfert galaxies selected from a mostly isotropic property, the flux at 60$\\mu$m, as well as for an additional 20 Seyfert galaxies with extended radio emission. The isophotes of the I band images were fitted with ellipses to determine the surface brightness profiles, the ellipticities and position angles of the host galaxy major axis. The parameters obtained with these fits were used to measure the surface brightness profiles in the B band. These images were also used to measure the integrated B and I magnitudes of the galaxies. These measurements will be combined with information from radio observations to study the orientation of radio jets relative to the host galaxy disk (Kinney et al. 2000)."
    },
    "0002/astro-ph0002242_arXiv.txt": {
        "abstract": "We use high precision multiband photometric data of the first overtone RR Lyrae \\uc to investigate the predictive capability of full amplitude, nonlinear, convective hydrodynamical models. The main outcome of this investigation is that theoretical predictions properly account for the luminosity variations along a full pulsation cycle. Moreover, we find that this approach, due to the strong dependence of this observable and of the pulsation period on stellar parameters, supply tight constraints on stellar mass, effective temperature, and distance modulus. Pulsational estimates of these parameters appear in good agreement with empirical ones. Finally, the occurrence of a well-defined bump just before the luminosity maximum gave the unique opportunity to calibrate the turbulent convection model adopted for handling the coupling between pulsation and convection. ",
        "introduction": "Variable stars play a key role in many astrophysical problems, since their pulsation properties do depend on stellar parameters, and therefore they can supply valuable and independent constraints on a large amount of current evolutionary predictions. In particular, the empirical evidence found long time ago in Magellanic Cepheids of the correlation between period and luminosity was the initial step for a paramount theoretical and observational effort aimed at using variable stars as standard candles to estimate cosmic distances. The current literature is still hosting a vivid debate on the intrinsic accuracy of the Cepheid distance scale (Bono et al. 1999; Laney 2000) and on the use of RR Lyrae variables to evaluate the distance -and the age- of Galactic globulars (Caputo 1998; Gratton 1998, G98).  Theoretical insights into the problem of radial stellar pulsations came from the linearization of local conservation equations governing the dynamical instability of stellar envelopes. Linear, nonadiabatic models typically supply accurate pulsation periods and plausible estimates (necessary conditions) on the modal stability of the lowest radial modes. However, a proper treatment of radial pulsations does require the solution of the full system of hydrodynamic equations, including a nonlocal and time-dependent treatment of turbulent convection (TC) to account for the coupling between radial and convective motions (Castor 1968; Stellingwerf 1982, S82).  The development of nonlinear, convective hydrocodes (S82; Gehmeyr 1992; Bono \\& Stellingwerf 1994, BS; Wuchterl \\& Feuchtinger 1998) gave the opportunity to provide plausible predictions on the properties of radial variables, and in particular on the topology of the instability strip, as well as on the time behavior of both light and radial velocity curves. This new theoretical scenario allowed to investigate, for the first time, the dependence of pulsation amplitudes and Fourier parameters on stellar mass, luminosity and effective temperature (see e.g. Kovacs \\& Kanbur 1997; Brocato et al. 1996; Feuchtinger 2000, F20). However, all these investigations dealt with parameters related to the light curve, whereas nonlinear computations supply much more information, as given by the detailed predictions of the light variation along a full pulsation cycle. Therefore, the direct comparison between observed and predicted light curves appears as a key test only partially exploited in the current literature (Wood, Arnold, \\& Sebo 1997, WAS).  In order to perform a detailed test of the predictive capability of our nonlinear, convective models we focused our attention on the photometric data collected by Heiser (1996, H96) for the field, first overtone -$RR_c$- variable \\uc. The reason for this choice relies on the detailed coverage of the U, B, and V light curves, as well as on the characteristic shape of the light curve, with a well-defined bump close to the luminosity maximum. This secondary feature provides a tight observational constraint to be nailed down by theory. Since the period of the variable strongly depends on the structural parameters (mass, luminosity, and radius) of the pulsator, the problem arises whether or not nonlinear pulsation models account for the occurrence of similar pulsators, and in affirmative how precisely the observed light curves can be reproduced by theoretical predictions.  In \\S 2 we present the comparison between theory and observations, while in \\S 3 we discuss the calibration of the TC model. Finally, in \\S 4 we briefly outline the observables which can further validate this theoretical scenario. ",
        "conclusions": "The comparison between theory and observations, namely the period and the shape of the {\\rm B} light curve of U Com, allowed us to supply tight constraints on the structural parameters such as stellar mass, effective temperature, and gravity, as well as on the distance of this variable. We also found that the occurrence of a well-defined bump close to the luminosity maximum can be safely adopted for constraining the metallicity of this object and for calibrating the TC model adopted for handling the coupling between convection and pulsation.  The approach adopted in this investigation seems quite promising since it only relies on nonlinear, convective models and on stellar atmosphere models. In fact, the best fit model to the empirical data was found by constructing sequences of iso-period models in which the stellar mass, the luminosity and the effective temperatures were not changed according to HB models but to the pulsation relation (BCCM). The comparison between theory and observations shows that both the structural parameters and the distance are in very good agreement with estimates available in the literature. No evidence for a systematic discrepancy was found in pulsation estimates, thus supporting the evidence that the individual fit to light curves can supply independent and firm constraints on the actual parameters and distances of variable stars. This finding confirms the results of a similar analysis on a LMC Bump Cepheid by WAS.  Accurate radial velocity data for \\uc are not available in the literature and therefore we could not constrain the accuracy of the velocity variation along the pulsation cycle. The three radial velocity points collected by FB98 agree quite well with the predicted curve. However, the radial velocity curve is a key observable for constraining the consistency of the adopted TC model (F20), and therefore new spectroscopic measurements of \\uc would be of great relevance for assessing the predictive impact of nonlinear, convective models. Theoretical observables of the best fit model discussed in this paper, as well as both radius and radial velocity variations are available upon request to the authors.   It is a pleasure to thank M. Groenewegen for providing us the \\uc distances based on the reduced parallax method and on the modified Lutz-Kelker correction, as well as for insightful discussions on their accuracy. We are indebted to R. Garrido for sending us radial velocity data and to T. Barnes for useful suggestions on current data. We also acknowledge an anonymous referee for some useful suggestions that improved the readability of the paper. This work was supported by MURST -Cofin98- under the project \"Stellar Evolution\". Partial support by ASI and CNAA is also acknowledged."
    },
    "0002/astro-ph0002074_arXiv.txt": {
        "abstract": "We measure simultaneously the properties of the energy spectra and the frequencies of the kilohertz quasi-periodic oscillations (QPOs) in fifteen low mass X-ray binaries covering a wide range of X-ray luminosities. In each source the QPO frequencies cover the same range of approximately $300$ Hz to $1300$ Hz, though the sources differ by two orders of magnitude in their X-ray luminosities (as measured from the unabsorbed 2--50 keV flux). So the X-ray luminosity does not uniquely determine the QPO frequency. This is difficult to understand since the evidence from individual sources indicates that the frequency and luminosity are very well correlated at least over short timescales. Perhaps beaming effects or bolometric corrections change the observed luminosities, or perhaps part of the energy in mass accretion is used to power outflows reducing the energy emitted in X-rays. It is also possible that the parameters of a QPO model are tuned in such a way that the same range of frequencies appears in all sources. Different modes of accretion may be involved for example (disk and radial) or multiple parameters may conspire to yield the same frequencies. ",
        "introduction": "Many low mass X-ray binaries exhibit quasi-periodic oscillations (QPOs) in their persistent X-ray flux in the kilohertz range as revealed by the Rossi X-ray Timing Explorer (RXTE). There are currently 18 such sources with published results.  Generally two kilohertz QPOs are observed simultaneously from a given system. In all cases, the QPOs are separated in frequency by about 250 to 350 Hz. The QPOs vary over a wide range in frequency. In 4U~0614+09, for example, the higher frequency QPO has been measured at frequencies between $449\\pm20$ Hz and $1329\\pm4$ Hz \\citep{vanstraaten00}. For reviews and references see \\citet{vanderklis_rev98b} and http://www.astro.uva.nl/$^{\\sim}$ecford/qpos.html.  The low mass X-ray binaries (LMXBs) which exhibit QPOs come in a wide variety. Most are persistent sources, but some transients are known with kilohertz QPOs: 4U~1608-52 \\citep{berger96,mendez98a}, Aql~X-1 \\citep{zhang98a}, and XTE~J2123-058 \\citep{homan99,tomsick99}.  The two traditional classes of LMXBs, Z and atoll sources \\citep{hk89}, have very similar QPOs, though the QPOs in Z-sources tend to have larger widths and smaller rms fractions. The X-ray dipper 4U~1915-05 \\citep{boirin00} also has shown kilohertz QPOs.  In all these systems, the kilohertz QPO frequencies are very similar, even though the inferred mass accretion rates differ by orders of magnitude \\citep{vdk_rev97a,vdk_rev97b}.  Here we quantify these comparisons by considering the ensemble of sources. The main tool is a measurement of the X-ray luminosity in each system simultaneous with a determination of its kilohertz QPO frequencies. This approach is inspired by the strong correlation of QPO frequency and count rate in individual sources. This correlation is very strict on short time scales \\citep[e.g. 4U~1728-34;][]{strohmayer96}, though on longer timescales of days to weeks in some sources a single correlation no longer holds \\citep[e.g. 4U~0614+09,][4U~1608-52,]{ford97a,mendez99a}. The same correlations are present if one considers X-ray flux instead of count rate \\citep{ford97b, zhang98a}. The QPO frequencies are clearly influenced to some extent by the X-ray luminosity.  Correlations of luminosity and kilohertz QPO frequency provide a rather direct connection to QPO models. In most current models, the frequency of one of the QPOs is set by the orbital frequency of matter in the inner disk \\citep{mlp98a,lai98,sv99,ot99b}.  Higher QPO frequencies are the result of faster orbital frequencies which are in turn coupled to higher mass accretion rates.  In the following we present simultaneous measurements of kilohertz QPOs and energy spectra in LMXBs. Section~2 details the analysis procedure and results with special notes on each source. Section~3 discusses these results in context with current QPO models. ",
        "conclusions": "Within a given low-mass X-ray binary the frequency of the kilohertz QPOs, \\freq, is well correlated with the X-ray flux \\citep{ford97b, zhang98a} or count rate \\citep{strohmayer96,wijnands98d,mendez99a,mendez99b,mendez99c}, at least on the timescale of about a day. Considering all the binaries as a group, however, such a correlation does not hold. This is a very clear feature of Figure~\\ref{fig:freqlx}, where \\freq covers roughly the same range of frequencies for sources of widely different X-ray luminosities, \\lx.  All sources have maximum frequencies at roughly 1000 to 1300 Hz, a fact that \\citet{zss97} have used to argue that the maximum \\freq is set by the orbital frequency at the marginally stable orbit.  In addition to the similar maximum \\freq, all the sources have roughly the same minimum \\freq and slope of their \\freq--\\lx relation. This is the central mystery presented here. How is it that \\lx and \\freq are decoupled in the ensemble of systems?  This decoupling has an apparent analog within individual sources.  In a given system, \\freq and \\lx (or flux, or count rate) are uniquely correlated within single observations spanning less than roughly a day. Between observations more widely separated in time, however, the correlation shifts and parallel lines appear in the \\freq vs \\lx diagram similar to those in Figure~\\ref{fig:freqlx}. Note, though, that these parallel lines in individual sources covers a much narrower range; flux shifts are a factor of a few at most in individual sources. 4U~0614+09 was first seen to have such parallel lines \\citep{ford97a,ford97b}, and the same effect is observed in Aql~X-1 \\citep{zhang98a}, 4U~1608-52 \\citep{mendez99a}, 4U 1728-34 \\citep{mendez99b}, and 4U~1636-53 \\citep{mendez99c}. There is a similar effect in Z-sources, where \\freq is correlated to the position on the instantaneous Z-track in the X-ray color diagram \\citep[e.g.][]{wijnands98d, jonker00} while the tracks themselves shift around in intensity.  One possible solution to the mystery of decoupled \\lx and \\freq is that the parameters of the mechanism producing the QPOs are tuned in such a way that \\freq is the same in all systems. As an example consider the magnetospheric beat-frequency model. A simple version of the theory predicts that the QPO frequency is set by $\\dot{M}/B^2$, where \\mdot is the mass accretion rate and $B$ is the surface magnetic field strength \\citep{as85}. The frequencies could then be the same if $B$ scaled in such a way that $\\dot{M}/B^2$ is constant in all systems \\citep{wz97}. Such a connection between \\mdot and $B$ was suggested previously on other grounds \\citep{hk89,pl97}. Other parameters, such as the neutron star spin, mass or temperature, might be involved as well, though it is not clear how these would fit into a detailed model.   The observational data do suggest that \\mdot has a role in setting the QPO frequency. The correlations of \\freq and \\lx suggest this, in as much as \\lx and \\mdot are related (see below). The timing properties point to a similar conclusion as well. The Fourier power spectra often show a noise component, whose power decreases with frequency above a break frequency of roughly 10 Hz.  The break frequency is strongly correlated with $\\nu_{kHz}$ \\citep{fk98,vanstraaten00,reig00,disalvo00}. The fact that the break frequency is thought to be a good indicator of $\\dot{M}$ \\citep{vanderklis94}, suggests that the frequency of the kilohertz QPO is also correlated with \\mdot. Another timing signal is the QPO at 10--50 Hz \\citep[e.g.][]{vanderklis96,fk98,pbk99} which also correlates with \\freq. Thus there are several timing features, all correlated with one another \\citep[see also][]{wk99,pbk99}. In addition \\freq also depends strongly on the energy spectra, sometimes parameterized as the distance along a track in the X-ray color diagram \\citep[e.g.][]{ vanderklis96,wijnands98a,zhang98c,mendez99a,mendez99c,kaaret99b}. The implication is that a single parameter underlies these correlations, and that parameter is likely \\mdot.  If there is a connection between \\freq and \\mdot, one might also expect a correlation of \\freq and \\lx, since \\lx is some measure of \\mdot. Why then is the range of \\freq similar for very different \\lx in Figure~\\ref{fig:freqlx}? In the following we consider one logical possibility: that \\lx and \\mdot do {\\em not} track one another.  Perhaps \\lx is simply not a good indicator of the bolometric luminosity and in fact the bolometric luminosity is similar in all systems.  In principle \\lx could misrepresent the bolometric luminosity just due to the limited 2--25 keV energy range of the RXTE/PCA. It is unlikely however that this is a large effect, since Beppo-SAX measurements from 0.1--200 keV indicate that not much energy is radiated outside the PCA band for these sources and our spectral models are applicable \\citep{piraino99}. Of course there could also be strong emission in the unobserved extreme ultraviolet band.  If the emission is not isotropic, the measured \\lx will also be an inaccurate indicator of the total emission. Inclination effects are one possibility: the lower \\lx sources may be viewed at a higher (more edge-on) inclination making \\lx smaller. This effect is well known in the dipping X-ray systems where the inclination is extremely edge-on and \\lx is low \\citep{parmar86}. An added attraction of this scenario is that it may explain the fact that Z-sources are strong radio emitters while the atoll-sources are not \\citep{fender00}. In this scenario, the less inclined, higher \\lx, Z-sources show strong radio emission because the radio jet is beamed into the line of sight, while atoll-sources at higher inclination and lower \\lx, are usually not detected in the radio because the radio jet is more in the plane of the sky. This may not be the full story, however, since the beaming would have to be narrow and a search for effects of inclination in the X-ray spectra with EXOSAT uncovered no evidence that inclination is important \\citep{wsp88}.  A general problem with preserving the same \\mdot in all the systems while changing the observed \\lx through anisotropy or bolometric corrections is that, if all the sources had the same \\mdot, they should all show the same X-ray burst properties. They do not; the Z-sources, for example, hardly burst at all \\citep{lewin93}. In the low-\\lx sources, \\mdot is also likely low because the persistent emission is at least 10 times weaker than in the bursts, some of which are at the Eddington limit. Assuming the anisotropy is about the same in the burst and persistent emission, \\mdot in these sources is then likely lower than in the sources near the Eddington limit, such as the Z sources.  Outflows are another way to decouple \\lx and \\mdot, and are a well-known feature of X-ray binaries, as seen for example in the collimated radio jets \\citep{hjellming95,fender99b}. One might expect that the outflows in the low-\\lx systems are stronger than those in the high-\\lx systems to preserve a similar accreted rate in the various systems. Radio observations, however, suggest that the opposite is true; the atoll sources are less luminous in radio than the Z-sources \\citep{fender00}.  Another alternative is that part of the \\mdot may be ineffective in determining \\freq while not being lost from the system. This could happen if the mass accretion rate occurs in a two component flow, radially and through a disk \\citep[e.g.][]{gl79,flm89,wijnands96}. The accretion rate through the disk is primarily responsible for setting \\freq, while the radial flow does not affect \\freq but does change $L_x$ \\citep[see][]{kaaret98}. \\citet{mlp98a} suggest that the disk accretion rate is similar in all sources. Matter is `scooped off' into a radial flow at the magnetospheric radius, and this process is more efficient in the higher \\lx sources because the fields are stronger. Under this scenario, the QPOs at higher \\lx should have a much smaller rms fraction due to the addition of unmodulated flux. This represents a problem for this scenario since the rms fraction apparently does not decrease enough with $L_x$ \\citep{ford00}.  All of the above effects may act to decouple \\lx and \\mdot. As outlined above, though, no single effect can account for the decoupling and each has problems.  If \\lx and \\mdot are unrelated, \\mdot can set the frequency of the QPOs while \\lx assumes any value, as observed."
    },
    "0002/astro-ph0002238_arXiv.txt": {
        "abstract": "We study the prospects for extracting detailed statistical properties of the Sunyaev-Zel'dovich (SZ) effect associated with large scale structure using upcoming multifrequency CMB experiments. The greatest obstacle to detecting the large-angle signal is the confusion noise provided by the primary anisotropies themselves, and to a lesser degree galactic and extragalactic foregrounds.   We employ multifrequency subtraction techniques and the latest foregrounds models to determine the detection threshold for the Boomerang, MAP (several $\\mu$K) and Planck CMB (sub $\\mu$K) experiments. Calibrating a simplified biased-tracer model of the gas pressure off recent hydrodynamic simulations, we estimate the SZ power spectrum, skewness and bispectrum through analytic scalings and N-body simulations of the dark matter.  We show that the Planck satellite should be able to measure the SZ effect with sufficient precision to determine its power spectrum and higher order correlations, e.g. the skewness and bispectrum.  Planck should also be able to detect the cross correlation between the SZ and gravitational lensing effect in the CMB.  Detection of these effects will help determine the properties of the as yet undetected gas, including the manner in which the gas pressure traces the dark matter. ",
        "introduction": "It is by now well established that the precision measurements of the cosmic microwave background expected from upcoming experiments, especially MAP and Planck satellite missions, will provide a gold mine of information about the early universe and the fundamental cosmological parameters (e.g., \\cite{Junetal96} 1996). These experiments can in fact do so much more.  With all-sky maps across the wide range of uncharted frequencies from $20$GHz-$900$GHz, the secondary science from these missions will arguably be as interesting as the primary science.  In this paper, we examine the prospects for extracting the large-scale properties of the hot intergalactic gas from multifrequency observations of the CMB. Inverse-Compton scattering of CMB photons by hot gas, known as the Sunyaev-Zel'dovich (SZ; \\cite{SunZel80} 1980) effect, leaves a characteristic distortion in the spectrum of the CMB, which fluctuates in the sky with the gas density and temperature. In the Rayleigh-Jeans (RJ) regime, it produces a constant decrement and with only low frequency measurements, the much larger primary anisotropies in the CMB itself obscure the fluctuations on scales greater than a few arcminutes (e.g., \\cite{GolSpe99} 1999). The upscattering in frequency implies an increment at high frequencies and a null around $217$GHz.  This behavior provides a potential tool for the separation of SZ effect from other temperature anisotropy contributors.   Since both the SZ spectrum and the CMB spectrum are accurately known, one can expect that foreground removal techniques developed to isolate the primary anisotropies can be reversed to recover the SZ signal in the presence of noise from the primary anisotropies.    Galactic and extragalactic foregrounds will be more challenging to remove.  Here we use the latest foreground models from \\cite{Tegetal99} (1999) that takes into account the fact that imperfect correlations in the foregrounds between frequency channels inhibits our ability to remove them. Using foreground information together with the expected noise properties of individual experiments, one can determine the  minimal detectable signal in each experiment and the upper limit achievable in the absence of detection.  Experiments with sufficient signal-to-noise can extract precision measurements for the power spectrum and higher order statistics such as the skewness.  Ultimately, they can provide detailed maps of the large-angle SZ effect.   To assess the prospects for an actual detection, we must model the SZ signal itself. The SZ effect is now routinely imaged in massive galaxy clusters (e.g., \\cite{Caretal96} 1996; \\cite{Jonetal93} 1993), where the temperature of the scattering medium  can reach as high as 10 keV, producing temperature changes in the CMB of order 1 mK at RJ wavelengths.  The possibility for detection of massive clusters in CMB satellite data has already been discussed in several studies (e.g., \\cite{Aghetal96}, \\cite{HaeTeg96} 1996, \\cite{Poietal98} 1998). Here, however, we are interested in the SZ effect produced by large-scale structure in the general intergalactic medium (IGM) where the gas is expected to be at $\\lesssim 1$keV in mild overdensities, leading to CMB contributions in the $\\mu$K range.  It is now widely believed that at least $\\sim$ 50\\% of the present day baryons, when compared to the total baryon budget from big bang nucleosynthesis, are present in gas associated with hot large-scale structure which has remained undetected given its temperature and clustering properties (e.g., \\cite{Fuketal98} 1998; \\cite{CenOst99} 1999; \\cite{Pen99} 1999). Recently, \\cite{Schetal00} (2000) has provided a tentative detection of X-ray emission from a large-scale filament in one of the deep ROSAT PSPC fields; previous attempts yielding upper limits are described in \\cite{KulBoh99} (1999) and \\cite{BriHen95} (1995). These results are consistent with current predictions for the X-ray surface brightness based on numerical simulations (e.g., \\cite{Cenetal95} 1995). \\cite{Pen99} (1999) argued that non-gravitational heating of the gas to $\\sim 1$keV is required to evade bounds from the soft X-ray background. These results suggest that the X-ray emission from this gas may be detectable in the near future with wide-field observations with Chandra X-ray Observatory\\footnote{http://asc.harvard.edu} and X-ray Multiple Mirror Mission\\footnote{http://astro.estec.esa.nl/XMM}.  On the theoretical front, hydrodynamic simulations of the SZ effect continue to improve (\\cite{daS99} 1999; \\cite{Refetal99} 1999; \\cite{Seletal00} 2000).  As a consensus from these simulations of basic properties such as the opacity weighted gas temperature and average Compton distortion is still lacking, we will base our assessment of the detectability of the large-scale SZ effect on a simple parameterization of the effect, based on a gas pressure bias model (\\cite{Refetal99} 1999), crudely calibrated with the recent hydrodynamic simulations. We employ perturbation theory, non-linear scaling relations, and N-body simulations for the dark matter to assess the statistical properties of the signal. Properly calibrated, these techniques can complement hydrodynamic simulations by extending their dynamic range and sampling volume. Currently, they should simply be taken as order of magnitude estimates of the potential signal.  Throughout this paper, we will take an adiabatic cold dark matter (CDM) model as our fiducial cosmology.  We assume cosmological parameters $\\Omega_c=0.30$ for the cold dark matter density, $\\Omega_b=0.05$ for the baryon density, $\\Omega_\\Lambda=0.65$ for the cosmological constant, $h=0.65$ for the dimensionless Hubble constant and a COBE-normalized scale invariant spectrum of primordial fluctuations (\\cite{BunWhi97} 1997).  The layout of the paper is as follows. In \\S~\\ref{sec:cleaning}, we describe the foreground and primary anisotropy removal method and assess their efficacy for upcoming CMB experiments. In \\S~\\ref{sec:sz}, we detail the bias model for the SZ effect and calculate through perturbation theory, analytic approximations and numerical simulations, the low order statistics of the SZ effect: its power spectrum, skewness and bispectrum. In \\S~\\ref{sec:sn}, having estimated the noise and the signal, we assess the prospects for measuring these low order statistics in upcoming experiments. We conclude in \\S~\\ref{sec:discussion} with a discussion of our results.   \\begin{figure*} \\centerline{\\psfig{file=figure1.eps,width=6.5in,angle=0}} \\caption{Top: foreground contributions to temperature anisotropies $(\\Delta T/T)^2 = l(l+1)C_l/2\\pi$ from the various foregrounds (dust, free-free, synchrotron, radio and infrared point sources, and rotating dust) at three fiducial frequencies as labeled. The SZ signal (solid, unlabeled) is estimated with the simplified model of \\S~3. Bottom: residual foregrounds after multifrequency subtraction for Boomerang, MAP and Planck. The total includes detector noise and residual CMB. } \\label{fig:clean} \\end{figure*} ",
        "conclusions": "\\label{sec:discussion}  We have studied the prospects for extracting the statistical properties of the Sunyaev-Zel'dovich (SZ) effect associated with hot gas in large-scale structure using upcoming multifrequency CMB experiments. This gas currently remains undetected but may comprise a substantial fraction of the present day baryons. The SZ effect has a distinct spectral dependence with a null at a frequency of $\\sim$ 217 GHz compared with true temperature anisotropies. This frequency dependence is what allows for effective separation of the SZ contribution with multifrequency CMB measurements.  As examples, we have employed the frequency and noise specifications of the Boomerang, MAP, Planck experiments. The MAP satellite only covers frequencies at RJ part of the frequency spectrum.  Consequently, only Boomerang and Planck can take full advantage of multifrequency separation of the SZ and primary anisotropies.  We have evaluated the detection threshold for SZ power spectrum measurements (see Fig.~\\ref{fig:error}). Boomerang and MAP should provide limits on the degree scale fluctuations at the several $\\mu$K level in rms; Planck should be able to detect sub $\\mu$K signals.  The expected level of the SZ signal in the fiducial $\\Lambda$CDM model is still somewhat uncertain.  We have employed a simple bias model for the pressure fluctuations, roughly normalized to recent hydrodynamic simulations (\\cite{Refetal99} 1999; \\cite{Seletal00} 2000), and calculated the resulting signal using analytic scaling relations and particle-mesh dark matter simulations. As hydrodynamic simulations improve, these techniques can be extended with more sophisticated modeling of the bias.  They complement hydrodynamic simulations by extending the dynamic range and simulated volume, the latter being important for questions of sample variance.  Assuming this simplified model of the SZ signal, Planck should have signal-to-noise per multipole of order unity for $l < 1000$.  Although the recovered maps are then somewhat noisy, they are sufficient for precise determinations of low order statistics such as the SZ power spectrum, bispectrum and skewness (see Figs.~\\ref{fig:simultwopt}-\\ref{fig:szbispec}). The skewness in principle can be used to separate the pressure bias from the underlying amplitude of the density fluctuations. The full bispectrum contains significantly more information but will be difficult to extract in its entirety. Current methods for measuring the bispectrum, tested with the COBE data, have concentrated at measuring specific modes such as $l_1 = l_2 =l_3 = l$ (\\cite{Feretal98} 1998). More work will clearly be required, especially in understanding the systematic errors at a sufficient level, but the wealth of information potentially present in the bispectrum should motivate efforts.  Note however that the non-Gaussianity in the SZ signal is not very strong due to the fact that it is constructed from many independent pressure fluctuations along the line of sight.  As a consequence, we expect that signal-to-noise ratios can be estimated by Gaussian approximations, but that techniques that try to improve the SZ-primary separation based on non-Gaussianity (\\cite{Hobetal98} 1998; \\cite{AghFor99} 1999) may not be particularly effective for this signal.  We caution the reader that our oversimplification of the SZ signal can cause problems for a naive interpretation of future detections. For example, \\cite{Seletal00} (2000) find that the SZ power spectrum in their simulations is dominated by shot noise from the rare hot clusters not included in our modeling. Fortunately since these contributions are highly non-Gaussian, they can can readily be identified and removed.   At the very least, $X$-ray bright clusters can be externally identified and removed; this has been shown to substantially reduce the shot noise contribution (\\cite{KomKit99} 1999). The effect we are modeling should be understood as the signal in fields without such clusters.  Another means of separating the SZ signal from large-scale structure from that of massive clusters is to cross correlate it with other tracers of large-scale structure that are less sensitive to highly overdense regions. An added benefit is that such a cross-correlation will also empirically measure the extent to which pressure fluctuations follow mass fluctuations. The CMB anisotropies themselves carry one such tracer in the form of the convergence from weak lensing.  It manifests itself as a three-point correlation or bispectrum (\\cite{GolSpe99} 1999) but without frequency information it is severely sample-variance limited due to confusion noise from primary anisotropies. Measuring the SZ-lensing correlation using the cleaned SZ maps improves the signal-to-noise for the detection by over an order of magnitude at degree scales.   Furthermore, the techniques introduced by \\cite{ZalSel99} (1999) provide a concrete algorithm for extracting most of the three-point signal without recourse to measuring all the configurations of the bispectrum. Conversely, SZ removal from the CMB maps themselves can assist in the detection of other smaller bispectrum signals by eliminating one source of confusion noise.   The cross-correlation coefficient between the SZ effect and CMB weak lensing is relatively modest ($\\sim$ 0.5, see \\cite{Seletal00} 2000). This is due to the fact that the SZ effect is a tracer of the nearby universe while CMB lensing is maximally sensitive to structure at $z\\sim 3$. A higher correlation is expected if SZ is cross-correlated with an external probe of low redshift structure. \\cite{PeiSpe00} (2000) suggested the cross-correlation of MAP CMB data and Sloan\\footnote{http://www.sdss.org} galaxy data. An improved approach would be to use the Planck derived SZ map rather than a CMB map. Using a SZ map reduces noise from the primary anisotropies and guarantees that any detection is due to correlations with the SZ effect. Extending the calculations in \\cite{PeiSpe00} (2000) with the Planck generated SZ map, we find signal-to-noise ratios which are on average greater by a factor of $\\sim$ 10 when compared to signal-to-noise values using MAP CMB map. In fact with redshifts for galaxies, Planck SZ map can be cross-correlated in redshifts bins to study the redshift evolution of the gas. Other promising possibilities include cross correlation with soft X-ray background measurements, as well as ultraviolet and soft X-ray absorption line studies.   All these considerations imply a bright future for SZ studies of the hot gas associated with large-scale structure with wide-field multifrequency CMB observations. Its detailed properties should be revealed in its non-Gaussianity and correlation with other tracers of large-scale structure."
    },
    "0002/astro-ph0002524_arXiv.txt": {
        "abstract": "Definitely, an affirmative answer to this question would have implications of fundamental importance for astrophysics (a new class of compact stars), and for the physics of strong interactions (deconfined phase of quark matter, and strange matter hypothesis). In the present work, we use observational data for the newly discovered millisecond X-ray  pulsar SAX J1808.4-3658 and for the atoll source 4U~1728-34 to constrain the radius of the underlying compact stars. Comparing the mass--radius relation of these two compact stars with theoretical models for both neutron stars and strange stars, we argue that a strange star model is more consistent with SAX J1808.4-3658 and 4U~1728-34, and suggest that they are likely strange star candidates. ",
        "introduction": "The possible existence of a new class of compact stars, which are made entirely of deconfined {\\it u,d,s} quark matter ({\\it strange quark  matter} (SQM)), is one of the most  intriguing aspects of modern astrophysics. These compact objects are called strange stars. They differ from neutron stars, where quarks are confined within neutrons, protons, and eventually within other hadrons ({\\it e.g. hyperons}). The investigation of such a possibility is relevant not only for astrophysics, but for high energy physics too. In fact, the search for a deconfined phase of quark matter is one of the main goals in heavy ion physics. Experiments at Brookhaven National Lab's Relativistic Heavy Ion Collider (RHIC) and at CERN's Large Hadron Collider (LHC), will hopefully clarify this issue in the near future.  The possibility that strange stars do exist is based on the so called {\\it strange matter hypothesis}, formulated by Witten~(1984) (see also Bodmer, 1971). According to this hypothesis, strange quark matter, in equilibrium with respect to the weak interactions, could be the true ground state of strongly interacting matter rather than $^{56}Fe$, {\\it i.e.} the energy per baryon of SQM must fulfil the inequality \\be \\bigg( {{E}\\over{A}} \\bigg)_{SQM} \\leq {{E(^{56}Fe)}\\over{56}} \\simeq 930~MeV, \\label{eq:stableSQM} \\ee at the baryon density where the pressure is equal to zero.  If the strange matter hypothesis is true, then a nucleus with A nucleons, could in principle lower its energy by converting to a strangelet (a drop of SQM).  However, this process requires a very high-order simultaneous weak interactions to convert about a number A of {\\it u} and {\\it d} quarks of the nucleus into strange quarks. The probability for such a process is extremely low {\\footnote{~It is proportional to $G_F^{2A}$, being $G_F$ the Fermi constant, and assuming a number $A$ of simultaneous weak processes.}}, and the mean life time for an atomic nucleus to decay to a strangelet is much higher than the age of the Universe. On the other hand, a step by step production of {\\it s} quarks, at different times, will produce hyperons in the nucleus, {\\it i.e.} a system (hypernucleus) with a higher energy per baryon with respect to the original nucleus. In addition, finite size effects (surface and shell effects) place a lower limit (A $\\sim$ 10--100) on the baryon number of a stable strangelet even if bulk SQM is stable (Farhi \\& Jaffe, 1984). Thus, according to the strange matter hypothesis, the ordinary state of matter, in which quarks are confined within hadrons, is a metastable state.  The success of traditional nuclear physics, in explaining an astonishing amount of experimental data, provides a clear indication that quarks in a nucleus are confined within protons and neutrons.  Thus, the energy per baryon $(E/A)_{ud}$ of {\\it u,d} quark matter (nonstrange quark matter) must be higher than the energy per baryon of nuclei \\be \\bigg( {{E}\\over{A}} \\bigg)_{ud} \\geq  930~MeV + \\Delta , \\label{eq:stableNSQM} \\ee being $\\Delta \\sim 4$~MeV a quantity which accounts for the lower energy per baryon of a finite chunk ($A \\sim 250$) of nonstrange quark matter with respect to the bulk ($A \\rightarrow \\infty$) case (Farhi \\& Jaffe, 1984). These stability conditions (eq.s (1) and (2)) in turn may be used to constrain the parameters entering in models for the equation of state (EOS) of SQM. As we show below, the existence of strange stars is allowable within the uncertainties inherent in perturbative Quantum Chromo-Dynamics (QCD). Thus {\\it strange stars may exist in the Universe}. ",
        "conclusions": ""
    },
    "0002/astro-ph0002181_arXiv.txt": {
        "abstract": "We present statistics of SGR 1806-20 bursts, combining 290 events detected with RXTE/PCA, 111 events detected with BATSE and 134 events detected with ICE. We find that the fluence distribution of bursts observed with each instrument are well described by power laws with indices 1.43, 1.76 and 1.67, respectively. The distribution of time intervals between successive bursts from SGR 1806-20 is described by a lognormal function with a peak at 103 s. There is no correlation between the burst intensity and either the waiting times till the next burst or the time elapsed since the previous burst. In all these statistical properties, SGR 1806-20 bursts resemble a self-organized critical system, similar to earthquakes and solar flares. Our results thus support the hypothesis that the energy source for SGR bursts is crustquakes due to the evolving, strong magnetic field of the neutron star, rather than any accretion or nuclear power. ",
        "introduction": "Soft gamma repeaters (SGR) are a rare class of objects characterized by their repetitive emission of low energy gamma-ray bursts. SGR bursts last $\\sim$ 0.1 s and their spectra are usually well described by an optically thin thermal bremsstrahlung (OTTB) model with kT $\\sim$ 20--40 keV. Three of the four known SGRs are associated with slowly rotating (P$_{\\rm spin}$ $\\sim$ ~5--8 s; Mazets et al. 1979, Kouveliotou et al. 1998, Hurley et al. 1999), ultra-strongly magnetized ($B \\gtrsim 10^{14}$ Gauss; Kouveliotou et al. 1998, Kouveliotou et al. 1999a) neutron stars positioned within or near young supernova remnants. For a review of the burst and persistent emission properties of SGRs, see Kouveliotou (1999b) and Hurley (2000).  Cheng et al. (1996) \\markcite{cheng96} reported similarities between particular statistical properties of a sample of 111 SGR 1806-20 bursts (observed with the International Cometary Explorer, ICE, between 1979 and 1984) and earthquakes. They noted that the distribution of the event energies of both phenomena follow a power law, dN $\\propto$ E$^{-\\gamma}$~dE, with index, $\\gamma$ $\\sim$ 1.6. Furthermore, they found that the cumulative waiting times between successive SGR bursts and earthquakes are similar. Laros et al. (1987) noted that the distribution of waiting times between successive SGR 1806-20 bursts follow a lognormal function, which was also seen between micro-glitches of the Vela pulsar (Hurley et al. 1994). Using the same data set, Palmer (1999) showed that, similar to earthquakes, some SGR 1806-20 bursts may originate from relaxation systems. {G\\\"o\\u{g}\\\"u\\c{s}} et al. (1999) studied a set of 1024 bursts from SGR 1900+14; 187 bursts were detected with the Burst and Transient Source Experiment (BATSE) aboard the Compton Gamma Ray Observatory (CGRO) and 837 bursts were detected with the Proportional Counter Array (PCA) on the Rossi X-ray Timing Explorer (RXTE) during an active period of the source in 1998. We found that their fluence distribution is consistent with a power law of index $\\gamma$ = 1.66 over 4 orders of magnitude. The distribution of waiting times between successive bursts also follows a lognormal function, which peaks at $\\sim$ 49 s. We discussed the idea that SGRs, like earthquakes and solar flares, are manifestations of self-organized critical systems (Bak, Tang \\& Wiesenfeld 1988). All of these results are consistent with the idea that SGR bursts are caused by starquakes, which are the result of a fracture of the crust of a magnetically-powered neutron star, or ``magnetar\" (Duncan \\& Thompson 1992\\markcite{dt92}; Thompson and Duncan 1995\\markcite{td95}, 1996\\markcite{td96}).  SGR 1806-20 exhibited sporadic bursting activity from the launch of BATSE (in April 1991) until November 1993 (Kouveliotou et al. ~1994\\markcite{kou94}). In October 1996, the source entered a burst active phase. The reactivation initiated a series of pointed observations with the RXTE/PCA over a period of two weeks. These observations led to the discovery of 7.47 s pulsations from SGR 1806-20 and confirmed its nature as a magnetar (Kouveliotou et al. 1998). In these two weeks RXTE/PCA recorded a total of 290 bursts\\setcounter{footnote}{6} \\footnote{Examples of  RXTE/PCA observations of SGR 1806-20 can be seen at {\\tt http://gammaray.msfc.nasa.gov/batse/sgr/sgr1806/}}. In the BATSE data, SGR 1806-20 burst activity was persistent but variable from October 1996 up to October 1999 with a total of 116 recorded bursts. In this {\\it Letter}, we present a comprehensive study of the statistical properties of SGR 1806-20 by combining several data bases. Sections 2, 3 and 4 describe the CGRO/BATSE, RXTE/PCA and ICE observations, respectively. Our results are presented in Section 5 and discussed in Section 6. ",
        "conclusions": "The fluence distributions of the SGR 1806-20 bursts seen with ICE and BATSE are well described by single power laws with indices 1.67 $\\pm$ 0.15 and 1.76 $\\pm$ 0.17, respectively, while RXTE bursts have an index of 1.43 $\\pm$ 0.06. These indices are similar to those found for SGR 1900+14 (1.66, {G\\\"o\\u{g}\\\"u\\c{s}} et al. 1999) and SGR 1627-41 (1.62, Woods et al. 1999b). The ICE and BATSE values are consistent with one another, over nearly the same energy range but at different epochs. This suggests that SGR event fluence distributions may not vary greatly in time, therefore, we combine the ICE and BATSE values to calculate a ``high-energy'' index, $\\gamma$ = 1.71 $\\pm$ 0.11. The difference between the ``low-energy'' (RXTE) index and the ``high-energy'' index is insignificant ($\\sim$ 2.3 $\\sigma$); more ``high-energy'' data are needed to determine whether there is a break in the distribution.  Power law energy distributions have also been found for earthquakes with $\\gamma$ = 1.4 to 1.8 (Gutenberg \\& Richter 1956; Chen et al. 1991; Lay \\& Wallace 1995), and solar flares, $\\gamma$ = 1.53 to 1.73 (Crosby et al. 1993, Lu et al. 1993). This is a typical behavior seen in self-organized critical systems. The concept of self-organized criticality (Bak, Tang \\& Wiesenfeld 1988) states that sub-systems self-organize due to some driving force to a critical state at which a slight perturbation can cause a chain reaction of any size within the system. SGR power law fluence distributions, along with a lognormal waiting time distribution support the idea that systems responsible for SGR bursts are in a state of self-organized criticality. We believe that in SGRs, the critical systems are neutron star crusts strained by evolving magnetic stresses (cf. Thompson \\& Duncan ~1995).  Cheng et al. (1996) suggested that there is a high energy cut-off in the cumulative energy distribution of SGR 1806-20 bursts seen by ICE. In a cumulative energy distribution, the values of neighboring points are correlated, consequently, judging the significance of apparent deviations is very difficult. For these reasons we used a maximum likelihood fitting technique and displayed the differential energy distributions (e.g Fig.2). We find no evidence for a high-energy cut-off in the ICE data of SGR 1806-20 up to burst energies $\\sim 10^{41}$ ergs. It should be noted, however, that a high energy cut-off or turnover must exist because otherwise the total energy diverges.  The distribution of waiting times of SGR 1806-20 bursts observed with RXTE is well described by a lognormal function, similar to that found by Hurley et al. (1994) for the bursts seen with ICE. The waiting times of the RXTE events are on average shorter than the ones observed with ICE, maybe due to different burst active phase of the source or to instrumental sensitivity (the PCA is more sensitive to weaker bursts than ICE, and the system displayed plenty of weaker bursts as well as strong ones in 1996), or combination of both. Recently  G\\\"o\\u{g}\\\"u\\c{s} et al. (1999) showed that the recurrence time distribution of SGR 1900+14 bursts observed with RXTE is also a lognormal function which peaks at $\\sim$ 49 s. The lack of any correlation between the intensity and the waiting time until the next burst agrees well with the results of ICE observations of SGR 1806-20 (Laros et al. 1987). This behavior, also seen in SGR 1900+14 (G\\\"o\\u{g}\\\"u\\c{s} et al. 1999) confirms that the physical mechanism responsible for SGR bursts is different from systems where accretion-powered outbursts take place (e.g. the Rapid Burster, Lewin et al. 1976, and the Bursting Pulsar, Kouveliotou et al. 1996)  The burst activity of SGR 1806-20 over the last three years is considerably different from that of SGR 1900+14. After a long period with almost no bursts, BATSE recorded 200 bursts from SGR 1900+14 between 1998 May and 1999 January, with a remarkably low activity thereafter. On the other hand, after SGR 1806-20 reactivated in 1996, it continued bursting on a lower rate, with 18 bursts in 1997, 32 in 1998 and 18 in 1999 through October. The latest RXTE observations of SGR 1806-20 in 1999 August revealed that smaller scale bursts are still occurring occasionally in this system, whereas contemporaneous RXTE observations of SGR 1900+14 do not show burst activity of any size. This continuation of burst activity may prevent the deposition of very large amounts of stress in the crust. Therefore, in SGR 1806-20 it may be less likely to expect, in the near future, a giant flare from this source, as the ones seen on 1979 March 5 from SGR 0526-66 (Mazetz et al. 1979) and on 1998 August 27 from SGR 1900+14 (Hurley et al. 1999)."
    },
    "0002/astro-ph0002148_arXiv.txt": {
        "abstract": "It is shown how wavelets may be used to analyse the absorption properties of the \\Lya forest. The Discrete Wavelet Transform of a QSO spectrum is used to decompose the light fluctuations that comprise the forest into orthogonal wavelets. It is demonstrated that most of the signal is carried by the moderate to lower frequency wavelets in high resolution spectra, and that a statistically acceptable description of even high signal--to--noise spectra is provided by only a fraction (10--30\\%) of the wavelets. The distributions of the wavelet coefficients provide a statistical basis for discriminating between different models of the \\Lya forest. The method is illustrated using the measured spectrum of Q1937--1009. The procedure described is readily automated and may be used to process both measured spectra and the large number of spectra generated by numerical simulations, permitting a fair comparison between the two. ",
        "introduction": "Measurements of QSO spectra show that the Intergalactic Medium (IGM) is composed of highly inhomogeneous structures. Ever since their identification by Lynds (1971) and the pioneering survey of Sargent \\etal (1980), these inhomogeneities have been described as discrete absorption systems, the \\Lya forest. With the view that the systems arise from individual intervening gas clouds, the \\Lya forest has been characterized using traditional absorption line statistics, most notably the line equivalent widths and, as the spectra improved in resolution and signal--to--noise ratio, the Doppler widths and \\HI column densities through Voigt profile line fitting to the features.  In the past few years, numerical simulations have successfully modelled many of the measured properties of the forest, showing that the absorption systems may arise as a consequence of cosmological structure formation (Cen \\etal 1994; Zhang, Anninos \\& Norman 1995; Hernquist \\etal 1996; Bond \\& Wadsley 1997; Zhang \\etal 1997; Theuns, Leonard \\& Efstathiou 1998). The simulations have shown, contrary to the picture in which the systems are isolated intergalactic gas clouds, that most of the systems originate in an interconnected web of sheets and filaments of gas and dark matter (Cen \\etal 1994; Bond \\& Wadsley 1997; Zhang \\etal 1998). Alternative statistical methods were subsequently introduced for describing the forest using the more direct measurements of the induced light fluctuations. These include the 1-point distribution of the fluctuations (Miralda-Escud\\'e \\etal 1996; Zhang \\etal 1997), and a quantity related to the 2-point distribution based on a weighted difference of the light fluctuations in neighbouring wavelength pixels (Miralda-Escud\\'e \\etal). A direct estimate of the 2-point transmission correlation function was made by Zuo \\& Bond (1994).  While the newer methods for analysing the \\Lya forest avoid the identification of absorption lines and the fitting of Voigt profiles, they are not necessarily fundamentally different in their description of the spectra. For instance, Zhang \\etal (1998) find that the distribution of optical depth per pixel in their simulation may be recovered by modelling the spectra entirely by discrete absorption lines with Voigt profiles. Rather the more direct methods circumvent a difficulty that has long plagued attempts to characterize the absorbers in terms of Voigt profiles: the sensitivity of the resulting line statistics to noise and to the fitting procedure. Absorption line fitting of necessity requires arbitrary decisions to be made regarding the setting of the continuum level, the deblending of features, and a decision on the acceptability of a fit. Different observational groups report different distributions for the line parameters. Most discrepant has been the inferred distribution of line widths. Even with the highest quality data gathered to date using the Keck HIRES, agreement is still lacking, with Hu \\etal (1995) finding a narrower Doppler parameter distribution with a significantly higher mean than found by Kirkman \\& Tytler (1997). The differences are important, as cosmological simulations predict comparable differences for a range of plausible cosmological models (Machacek \\etal 2000; Meiksin \\etal 2000).  The purpose in this paper is to develop a method that provides an alternative objective description of the statistics of the \\Lya forest. Ultimately the goal is to employ the same method for analysing both observational data and data derived from numerical simulations in order to compare the two on a fair basis. Because of the large number of synthetic spectra generated from a simulation necessary to provide a correct average description of the forest, two principal requirements of the procedure are that it be fast and easily automated. Although automated or semi--automated Voigt profile fitting procedures exist (AutoVP, Dav\\'e \\etal 1997; VPFIT, developed by Carswell and collaborators), these procedures still require arbitrary decisions to be made to obtain acceptable fits. The complexity of the codes makes it difficult to assess the statistical significance of differences between the measured distributions of the absorption line parameters and those predicted. The codes also are computationally expensive, making very costly their application to the large number of simulated spectra required to obtain a statistically valid average of the line parameters. For these reasons, a faster less complex method would be desirable. The Voigt profile fitting codes yield important parameters, like the linewidths, which contain physical information (eg, gas temperature and turbulent velocities), that the direct-analysis methods do not. It would thus be desirable for an alternative method to retain some of this information. The method presented here utilizes wavelets to characterize the absorption statistics of the \\Lya forest. It is not intended to be a replacement for Voigt profile fitting, but a fast alternative that allows a ready comparison between the predictions of numerical models and measured spectra and a clear statistical analysis of the results.  The outline of the paper is as follows: in \\S\\ref{sec:wavelets} it is shown how the statistics of the \\Lya forest may be characterized using wavelets. In \\S\\ref{sec:results} the method is applied to the measured spectrum of a high redshift QSO. The results are summarized in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  Wavelets may be usefully employed to provide a statistical characterizaton of the absorption properties of the \\Lya forest. An approach is presented that performs a multiresolution analysis of the forest using the Discrete Wavelet Transform of the QSO spectrum. The transform decomposes the local frequency dependence of the light fluctuations into an orthogonal hierarchy of basis functions, the wavelets. It is found that in spectra of better than 10~\\kms resolution, most of the information of the spectrum is carried by the lower frequency wavelets. For a signal--to--noise ratio typical of even the highest quality spectra ($S/N=10-100$), only 10--30\\% of the wavelets are required to provide a statistically acceptable description of the spectrum, corresponding to a data compression factor of 3--10. It is shown that a Voigt profile line analysis performed on the wavelet filtered spectra yields nearly identical line parameter distributions as obtained from the original unfiltered spectra.  The distributions of the wavelet coefficients offer an alternative statistical description of the \\Lya forest while retaining information on the line widths. It is demonstrated that the correlations of coefficients between different levels in the wavelet hierarchy are weak (a few percent or smaller). Consequently, each of the distributions may be treated as statistically independent to good approximation.  The method is applied to a Keck HIRES spectrum of Q1937--1009. The wavelet coefficient distributions behave qualitatively similarly to those found in Monte Carlo simulations based on the line parameter distributions reported by Kirkman \\& Tytler. The measured distributions, however, show some differences on the scale $17-34$~\\kms.   The results demonstrate that Multiresolution Analysis using the Discrete Wavelet Transform provides an alternative objective, easily automated procedure for analysing the \\Lya forest suitable for basing a comparison between the measured properties of the \\Lya forest and the predictions of numerical models.  \\bigskip The author thanks S. Burles and D. Tytler for kindly providing the spectrum of Q1937--1009, and R. Dav\\'e for permission to use AutoVP."
    },
    "0002/hep-ex0002048_arXiv.txt": {
        "abstract": "The Alpha Magnetic Spectrometer (AMS) was flown on the space shuttle {\\it Discovery} during flight STS--91 in a 51.7$^\\circ$ orbit at altitudes between 320 and 390\\,km. A total of $2.86\\times 10^6$ helium nuclei were observed in the rigidity range 1 to 140\\,GV. No antihelium nuclei were detected at any rigidity. An upper limit on the flux ratio of antihelium to helium of $< 1.1\\times 10^{-6}$ is obtained. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002476_arXiv.txt": {
        "abstract": "This paper presents some statistical correlations of 72 northern spiral galaxies. The results show that early-type spirals that are brighter, and thicker, and the axis ratios ($H_{z}/H_r$) of the disk tend to be smaller along the Hubble sequence. We also find that $H_{z}/H_r$ correlates strongly with the galaxy's color, and early-type spirals have larger values of $H_{z}/H_r$. The inclinations obtained by fitting the pattern of a spiral structure with a logarithmic spiral form are nearly the same as those obtained by using the formulas of Aaronson et al. (1980). Finally, the mean measured pitch angles for the different Hubble sequences in the Third Catalogue of Bright Galaxies by de Vaucouleurs et al. (1991) are derived. ",
        "introduction": "The inclination of a spiral galaxy ( i.e., the angle between the galactic plane and the tangent plane of the celestial sphere) is not only an important parameter, but also difficult to determine. A spiral galaxy consists of a thin disk, a bulge and spiral arms that are thought to be situated in the disk. If we assume that the thickness of the spiral plane is rather negligible in comparision to its extension, and that when a spiral galaxy is inclined moderately to the plane of sky, the thickness of the nucleus can be omitted, the inclination $(\\gamma)$ can be obtained by:  \\begin{equation} \\gamma=\\arccos(\\frac{d}{D}), \\end{equation} where $D$ and $d$ are the apparent major and minor isophotal diameters respectively. When a spiral galaxy is seen edge-on, it is not possible to consider the thickness of the nuclear part as negligible. Thus, Eq. (1) cannot be used to calculate the inclination. The reason for this is that the apparent minor isophotal-diameter consists of two parts. One is attributed by the disk and, another by the bulge, the latter of which will decrease the real value of the inclination.  Considering that the disk is not infinitely thin, Aaronson et al. (1980) corrected Eq. (1) by  \\begin{equation} \\gamma=\\arccos\\sqrt{1.042(\\frac{d}{D})^{2}-0.042}+3^{\\circ}. \\end{equation} The constant of $3^{\\circ}$ is added in accordance with an empirical recipe. A more elaborate specification of the axial ratio for an edge-on system that depends on the Hubble type could be justified. The thinnest galaxies are Sc spirals', earlier types have larger bulges. Giovanelli et al. (1997) provided an example to justify why they assumed the axial ratio of Sc galaxies to be 0.13. A smaller value of the axial ratio for an edge-on system results in smaller derived inclinations, where the spirals are more face-on. Besides, if the values $D$ and $d$ are approximations due to errors, the inclination obtained by Eq. (2) is not an exact value. Ma et al. (1997, 1998) proposed a method to determine the inclination of a spiral galaxy by fitting a spiral arm with a logarithmic spiral form with constant pitch angle. They obtained the inclinations of 72 northern spiral galaxies.  The question of the mathematical form of spiral arms was recognized at the beginning of this century (von der Pahlen, 1911; Groot, 1925). Then, Danver (1942), Kennicutt (1981) and Kennicutt \\& Hodge (1982) systematically studied the shapes of spiral arms. Using the method of the least squares and as many points as possible situated on the spiral arm in question, Danver (1942) studied a sample of 98 nearby spirals by drawing the projected images on white paper and then, by copying it on the paper to be used for the measurement thanks to transparent illumination. Kennicutt (1981) measured the shapes of spiral arms in 113 nearby Sa-Sc galaxies by disposing directly of photographic enlargement and using an iterative procedure to correct for inclination effects. He gave an initial estimate of the inclination and pitch angle to orient the spiral to a face-on geometry, and then used any residual sinusoidal deviations in the arm shapes to make small corrections to the derived orientation. Using the IRAF software, Ma et al. (1997, 1998) fitted the shapes of spiral arms on the images, so that they could show clearly whether the fitting was good or not. The DISPLAY program of IRAF software can enlarge the image and change its grey scale to minimize any personal prejudice about the regularity and prominence of arms. But we must emphasize that the DISPLAY program of IRAF has many variables, so the results are not always objective. In our program, we modify z1 (minimum greylevel to be displayed) and z2 (maximum greylevel to be displayed) in the DISPLAY program in order to display the images clearly. In the procedure of fitting, we emphasize the global spiral structure, where, except for the small-scale distortions, the arms can be represented by the logarithmic spiral forms.  There has been much interest concerning the separation of disk and bulge components in the observed surface brightness distribution of spiral galaxies. de Vaucouleurs (1958), for instance, established an isophotic map of M~31 in blue light by means of direct photoelectric scans, spaced at 10$'$ intervals in declination from +39$^{\\circ}$31$'$ to 42$^{\\circ}$30$'$. From photoelectric photometry, he determined that the thickness of the flat component is about 0.8 kpc. By assuming that a galaxy has an infinitesimally thin disk, Freeman (1970) and Sandage et al. (1970) collected and studied the radial distribution of the surface brightness $I(r)$ for thirty-six S0 and spiral galaxies, and showed that $I(r)$ distribution for these galaxies can be presented by two main components: an inner spheroidal component which follows the law of \\begin{equation} \\log I(r) \\propto r^{1/4} \\end{equation} and an outer exponential component (disk), with \\begin{equation} \\log I(r)=I_0e^{-r/h_r}, \\end{equation} where $h_r$ is defined as a radial scale length. Van der Kruit and Searle (1981a) proposed a model for the light distribution in the disks of edge-on spiral galaxies, assuming that a galaxy has a locally isothermal, self-gravitating and truncated exponential disk. This model has the feature of being isothermal in $z$ at all radii with a scale parameter $z_0$ and has an exponential dependence of surface brightness upon $r$ with a scale length $h_r$. The space-luminosity of this model can be described by \\begin{equation} L(r, z)=L_0e^{-r/h_r}{\\rm{sech}}^{2}(z/z_0). \\end{equation} With this model, van der Kruit \\& Searle (1981a, 1981b, 1982a, 1982b) determined $h_r$ and $z_0$ for seven disk-dominated and one spheroid-dominated spiral galaxies by using three-color surface photometry. Peng et al (1979) investigated three-dimensional disk galaxies, based on the fundamental assumption by Parenago that the density distribution along $z$-direction for a finite thickness disk is \\begin{equation} \\rho(r, \\phi, z)=\\frac{1}{H_z}\\sigma(r, \\phi)e^{-|z|/h}, \\end{equation} where $h$ is defined as an exponential scale height, $H_{z}$ is defined as a thickness of disk and equals $2h$, and $\\sigma(r, \\phi)$ is the surface density. By solving Poission's equation for a logarithmic density perturbation, Peng et al. (1979) obtained a criterion for density waves to appear, which is \\begin{equation} r>r_0=\\frac{H_z\\sqrt{m^2+\\Lambda^2}}{2}, \\end{equation} where $(r_0, \\phi_0)$ is the polar coordinate of the starting point from which arms of a galaxy stretch outward on the galactic plane, and $m$ is the number of the arms in a spiral galaxy. Based on this criterion, Peng (1988) proposed a method for estimating the thickness of a non-edge-on spiral, and derived the thicknesses of four galaxies.  Guthrie (1992) derived the axial ratios $R$ of disc components for 262 edge-on spiral galaxies on print copies of the blue Palomar Sky Survey plates by using a microscope fitted with a micrometer eyepiece. He then analyzed the distribution of isophotal axial ratios for 888 diameter-limited normal Sa-Sc galaxies to give information on the true axial ratios $R_0$, and at last presented the mean value of $\\log R_0$ is $0.95\\pm0.03$.  Ma et al. (1997, 1998) derived the thicknesses of 72 spirals by using Peng's proposal (Peng, 1988) and presented some statistical correlations between thickness or flatness and other parameters.  The value of $R_0$ derived by Guthrie (1992) is based on observational work and, should be much more reliable. So, it is important to compare Ma et al.'s results (1997, 1998) to Guthrie's (1992). In Ma et al. (1997, 1998), the flatnesses for 72 galactic disks ($H_{z}/D_0$) \\footnote{$D_0$, which is measured at or reduced down to the surface brightness level $\\mu_{B}=25.0 B$ magnitudes per square arcsecond, and corrected to the ``face-on'' ($\\gamma = 0^{\\circ}$). For the Galactic extinction, but not for redshift, $D_0$ is from the Third Catalogue of Bright Galaxies by de Vaucouleurs et al. (hereafter RC3).} were given, and the mean value is $0.033 \\pm 0.002$. Suppose that the value of ratio of radial scale length ($h_r$) over exponential scale height ($h$) for an average exponential galactic disk is equal to $R_0$ from Guthrie (1992), which is 9. From Freeman (1970), we can derive $D_0/(2h_r)\\approx 5$. Thus, we obtain $H_{z}/D_0\\approx 0.023$, which is in relative agreement with the mean value of Ma et al. (1997, 1998). At the same time, we calculate the values of $\\overline{\\log R_0}$ for spirals of various types T and list them in Table 1. T, from the RC3, is morphological types, and $\\sigma$ is the dispersion. From this table and Table 3 of Guthrie (1992), we can see that our results are in agreement with Guthrie's (1992). Except for the data concerning Scd galaxies, our results also agree with de Grijs' (1998).  \\begin{table} \\caption{Values of $\\overline{\\log R_0}$ for spirals of various types T.} \\begin{tabular} {c|ccc} \\hline T & $\\overline{\\log R_0}$ & $\\sigma$ \\\\ \\hline 2 & $0.65\\pm 0.04$ & 0.09\\\\ 3 & $0.79\\pm 0.06$ & 0.21\\\\ 4 & $0.79\\pm 0.04$ & 0.21\\\\ 5 & $0.90\\pm 0.04$ & 0.18\\\\ 6 & $1.20\\pm 0.15$ & 0.33\\\\ \\hline \\end{tabular} \\end{table}  The structure of this paper is as following: in Sect. 2, we outline the principles of obtaining an inclination and a pitch angle; Sect. 3 presents some statistical properties; and conclusions will be shown in Sect. 4. ",
        "conclusions": "In this paper, we investigate some statistical correlations about the thickness of galactic disks and pitch angles of spiral arms. The main conclusions are: (1). Early-type spirals, that are brighter on the average, are thicker; (2). The axis ratio ($H_{z}/H_r$, here $H_{r}$ and $H_z$ are defined as two times the radial scale length ($h_r$) and exponential scale height ($h$)) of galactic disk tends to be smaller along the Hubble sequence; (3). Except for a few galaxies, early-type spiral galaxies have larger values of $H_{z}/H_r$; (4). $H_{z}/H_r$ correlates strongly with galaxy color; (5). The inclinations obtained by fitting the pattern of spiral structure with a logarithmic spiral form are nearly the same as those obtained by using the formulas of Aaronson et al. (1980); (6). The mean measured pitch angles for different Hubble sequences in the RC3 are derived, and the results show that the mean pitch angles of Sa-Sc's are not larger than $15.5^{\\circ}$, so that cotan$\\mu\\geq3.6$. Thus the  WKB approximation can be satisfied at the mean pitch angles from Sa-Sc's, but not by very much; (7). From early to late Hubble types, the mean value of pitch angle increases, despite some scattering.  Although the method, proposed by Peng (1988) for deriving the thickness of a face-on disc, is effective and simple, it relies on a spiral structure theory that predicts that a spiral arm has to end somewhere in the disk. However, this theory might not be completely right. For example, Zaritsky et al. (1993) has presented K-band (2.2 $\\mu m$) images of M~51, which reveal remarkable dynamical structure not visible in the conventional optical observations, and show that the spiral arms extend significantly further towards the galaxy's center than previously observed. In the optical images, the spiral arms begins at a radius of about $30^{''}$ from the center. But, the K-band residual image showed the spiral arms wind through an additional $540^{\\circ}$ beyond that seen in the B-band images of the entire galaxy, and end at about $10^{''}$ from the center of the galaxy. If spiral galaxies whose spiral arms do not stop anywhere in the disk exist, the parameters $\\rho_0$ and $r_0$ (Peng, 1988; Ma et al., 1997, 1998) cannot be found. In the K-band images, which minimize the effect of dust and maximize sensitivity to the dominant stellar population, we can derive our reliable values of $\\rho_0$ and $r_0$. As an example, we apply our method to the image of M~51 in the K-band\\footnote{The image of M~51 in K-band is provided by Prof. Zaritsky.} and derive the flatness of disk ($H_z/D_0$), which is $0.010\\pm 26.4\\%$. Comparing this value with Peng's (1988) ($H_z/D_0=0.013\\pm 16.4\\%$), we can find that the value of the flatness based on the K-band image is smaller. The reason is that, in the K-band image, the effect of dust can be minimized and the values of $\\rho_0$ and $r_0$ may be reliably derived. We also emphasize that the images, which Ma et al. (1998) used, are from the Digitized Palomar Sky Survey, in which many images have burnt-out centers. There are some pictures that have burnt-out centers in Ma et al. (1998), but we can change the parameters of DISPLAY program to minimize the effect."
    },
    "0002/astro-ph0002195_arXiv.txt": {
        "abstract": "We present an inventory of mid-infrared spectral features detected in high resolution (R$\\sim$1500) ISO-SWS 2.4--45$\\mu$m spectra of the galaxies \\object{M\\,82}, \\object{NGC\\,253}, \\object{Circinus}, \\object{NGC\\,1068}, and a position in the \\object{30\\,Doradus} region of the Large Magellanic Cloud. We discuss their identifications and highlight possible relations between these features and the physical state of the interstellar medium in galaxies. The spectral features vary considerably from source to source in presence and relative strength. Emission features are largely absent in the intense radiation field close to an AGN. Compared to normal infrared-selected starbursts, they also seem to be weaker in a low metallicity, intensely star forming environment. The large number of features beyond 13$\\mu$m is remarkable. Some of the features have -- to our knowledge -- not been reported before in astronomical objects.  In the 5--13$\\mu$m region, emission from unidentified infrared bands (UIBs), usually ascribed to aromatic molecules, and apparent silicate absorption dominate the spectrum. The density of features makes it difficult to determine the continuum, particularly in ground-based data of limited wavelength coverage. In fact the apparent depth of the 9.7$\\mu$m silicate absorption may be overestimated in the presence of UIB emission, as we demonstrate by comparing the spectrum of M\\,82 to the (absorption free) spectrum of the reflection nebula \\object{NGC\\,7023}. No strong silicate absorption is present in M\\,82. The (very small grain) dust continuum under the UIB emission in our starburst templates can be modeled by a simple power law, starting at wavelengths between 8 and 9$\\mu$m.  We find broad H$_2$O-ice absorption features at 3.0$\\mu$m in M\\,82 and NGC\\,253. Their optical depths (relative to the visual extinction) indicate that the lines of sight towards these galaxies have similar properties as the line of sight towards the Galactic Center. The active galaxy NGC\\,1068 exhibits a clearly different spectrum of absorption features, indicating different physical conditions in the obscuring regions of this AGN compared to the starburst templates.   The spectra are valuable templates for future mid-infrared missions. We smooth our data to simulate low resolution spectra as obtained with ISOCAM-CVF, ISOPHOT-S, and in the future with the low resolution mode of SIRTF-IRS, and use our high spectral resolution information to highlight possible identification problems at low resolving power that are caused by coincidences of lines and features. The spectra are available in electronic form from the authors. ",
        "introduction": "Mid-infrared spectra of galaxies are rich in emission lines, and display prominent broader emission and absorption features due to the presence of various solids and/or large molecules in their interstellar medium (ISM). Significant variation from source to source suggests that these features may provide important diagnostics of the ISM conditions in galaxies.  The ground based and Kuiper Airborne Observatory spectra of the prototypical starburst M\\,82 by Gillett et al. (1975) and Willner et al. (1977) fully established the existence of the mid-infrared `unidentified infrared bands' (UIB) at 6.2, 7.7, 8.6, and 11.3$\\mu$m in galaxy spectra. These emission bands are characteristic of C-C and C-H bonds in aromatic molecules. In this paper we will refer to them as `PAH features' according to one of the most popular interpretations of their carrier, polycyclic aromatic hydrocarbon molecules \\footnote{Other suggested carriers include small grains of hydrogenated amorphous carbon (HACs), quenched carbonaceous composites (QCCs), or coal.}. These detections and related work using the IRAS LRS (Cohen \\& Volk 1989) form the basic pre-ISO knowledge of mid-infrared spectral features in galaxies.  Considerable work has also been done from the ground but has been limited to the features found in atmospheric windows, mainly silicate absorption and PAH feature emission in the N band (e.g. Roche et al. 1991, and references therein) and the companion PAH feature in the L band (e.g. Moorwood 1986). The restriction to atmospheric windows increases problems in establishing the `continuum' on which the features are superposed. This is a nontrivial task, even with full wavelength coverage, due to the crowding of mid-infrared emission and absorption features (especially in the 10$\\mu$m region).  With the Short Wavelength Spectrometer SWS (de~Graauw et al. 1996) on board the Infrared Space Observatory ISO (Kessler et al. 1996) high spectral resolution observations with good signal-to-noise (S/N) were obtained for a number of bright galaxies. Their main advantages lie in continuous wavelength coverage from 2.4 to 45$\\mu$m and in the possibility to clearly separate features from nearby emission lines.  The interpretation of galaxy-integrated spectra strongly benefits from comparisons to similar observations of galactic sources, sometimes spatially resolved, allowing better isolation of the physical mechanisms at work. Recent ISO spectra of many galactic template objects, such as reflection nebulae (e.g. Boulanger et al. 1996, Cesarsky et al. 1996a, Verstraete et al. 1996, Moutou et al. 1998), planetary nebulae and circumstellar regions (e.g. Beintema et al. 1996), and HII regions (Roelfsema et al. 1996, Cesarsky et al. 1996b) have clearly demonstrated the importance of such template spectra. They prove that PAHs are an ubiquitous part of the ISM. Additional information on emission features of crystalline silicates comes from similar template observations with ISO of e.g. planetary nebulae (Waters et al. 1998), evolved stars (Waters et al. 1996), young stars (Waelkens et al. 1996), or LBVs in the LMC (Voors et al. 1999). Absorption features (silicates, ices) have been found e.g. in the Galactic center (Lutz et al. 1996, Chiar et al. 2000),  young stellar objects (d'Hendecourt et al. 1996, Whittet et al. 1996, Dartois et al. 1999), and in dark clouds in the solar neighborhood (Whittet et al. 1998).  In this paper we present an inventory of mid-infrared spectral features detected in high resolution (R$\\sim$1500) ISO-SWS 2.4--45$\\mu$m spectra of the starburst galaxies M\\,82 and NGC\\,253, the Seyfert 2 galaxies Circinus and NGC\\,1068, and a position in the 30\\,Doradus star forming region of the Large Magellanic Cloud (Sect. \\ref{s:inventory}). We briefly discuss possible feature identifications (Sect. \\ref{s:ident}) and highlight possible relations between these features and the physical state of the interstellar medium in galaxies (Sect. \\ref{s:PAH_var}). We also address the issue of the continuum determination and the apparent depth of the silicate absorption at 9.7$\\mu$m (Sect. \\ref{s:continuum}). Finally (Sects. \\ref{s:lowres} and \\ref{s:conclusions}) we demonstrate the use of these ISO spectra as templates for future infrared missions such as SIRTF, with particular emphasis on potential identification problems at low resolving power that are caused by coincidences of lines and features.  All the spectra shown here exhibit a large number of atomic, ionic and molecular emission lines. These have been or will be discussed elsewhere, along with more details on observations and data processing (Circinus: Moorwood et al. 1996; M\\,82: Lutz et al. 1998b, Schreiber 1998; NGC\\,1068: Lutz et al. 2000; 30\\,Dor: Thornley et al., in prep.). ",
        "conclusions": "\\label{s:conclusions}  We have detected a large number of mid-infrared features in galaxy spectra, some of them previously unobserved, and discussed the dependence of the dust features on ISM condition in galaxies. The spectral features vary considerably from source to source in presence and relative strength. Emission features are largely absent in the intense radiation field close to an AGN, and weak in a low metallicity, intensely star forming environment. Differences in the absorption spectra point to different physical properties of the obscuring regions in starburst and active galaxies.  The spectra presented here will be valuable template spectra for future mid- and far-infrared space missions such as SIRTF, SOFIA or FIRST. They provide important clues for the identification and interpretation of high redshift, dusty galaxies. The strongest PAH features can be used to provide redshift information in far-infrared photometric galaxy surveys (Simpson \\& Eisenhardt 1999, see also the example of 21396+3623, Rigopoulou et al. 1999). Furthermore, they affect galaxy number counts. For instance, Xu et al. (1998) have constructed semi-empirical galaxy SEDs to model the considerable PAH effects on number counts and redshift distributions. Finally, the continuum and the PAH features can be used to distinguish between starburst activity and active nuclei in high redshift galaxies, as has been demonstrated for local infrared bright galaxies (Genzel et al. 1998, Lutz et al. 1998a, Rigopoulou et al. 1999).  The advantage of the wide wavelength coverage of the SWS spectra has been used to illustrate the problem of the continuum definition and the true depth of the silicate absorption. We find that in our starburst templates the hot VSG dust continuum begins to rise around 8 to 9$\\mu$m, and that it can be well fitted by a simple power-law up to 20...25$\\mu$m. Finally we have demonstrated possible line identification problems in low resolution spectra.  The spectra presented here are available in electronic form from the authors. We want to note again, that different parts of the spectra were observed through different aperture sizes, which should be taken into account for a detailed use as template spectra."
    },
    "0002/astro-ph0002530_arXiv.txt": {
        "abstract": "Deep I-band imaging to I~$\\approx$~26.5 of the soft gamma--ray repeater SGR 1900+14 region has revealed a compact cluster of massive stars located only a few arcseconds from the fading radio source thought to be the location of the SGR \\citep{fra99}. This cluster was previously hidden in the glare of the pair of M5 supergiant stars (whose light was removed by PSF subtraction) proposed by \\citet{vrb96} as likely associated with the SGR 1900+14. The cluster has at least 13 members within a cluster radius of $\\approx$~0.6~pc based on an estimated distance of 12--15 kpc. It is remarkably similar to a cluster found associated with SGR 1806--20 \\citep{fuc99}. That similar clusters have now been found at or near the positions of the two best--studied SGRs suggests that young neutron stars, thought to be responsible for the SGR phenomenon, have their origins in proximate compact clusters of massive stars. ",
        "introduction": "The \\citet[V96]{har96,vrb96} survey of the original Network Synthesis Localization (NSL) of SGR 1900+14 \\citep{hur94} found a pair of nearly identical M5 supergiant stars, separated by 3.3 arcsec, and at an estimated distance of 12-15 kpc. While just outside of the original NSL, they lie within the ROSAT HRI localization of the quiescent X--ray source RX~J190717+0919.3 thought to be associated with SGR 1900+14 \\citep{hur96}. On the basis of the small probability that even one supergiant would lie within the ROSAT error circle and that at least one other supergiant had been associated with an SGR (1806--20; \\citet{van95,kul95}), V96 proposed that the M star pair may be associated with the SGR 1900+14 source. The position of the M star pair has continued to be consistent with more recent X--ray and gamma--ray observations which, taken together, have narrowed considerably the actual location of SGR 1900+14 from the original NSL area of 5 arcmin$^2$.  These recent X--ray and gamma--ray observations have also detected variations with a period of 5.16 sec \\citep{hur99a,mur99,kou99} and a deceleration of $\\dot P \\approx 10^{-10}$~sec/sec.  Taken together, these are interpreted as evidence that the SGR source is a magnetar, though there remains some uncertainty in this interpretation \\citep{mar99}.  Additionally, a variable and fading radio source was detected shortly after the 27 August SGR 1900+14 superburst by \\citet{fra99}, providing strong evidence that it was the radio counterpart to the SGR. Its subarcsec accurate position is located only a few arcseconds from the M stars. These positional coincidences, the lack of a plerionic radio source, and, despite arguments for SNR G42.8+0.6 in the literature, the lack of a coincident supernova remnant, suggest that the system of proximate, high--mass M stars should not yet be dismissed as an evolutionary companion to the pulsating X--ray source associated with the SGR.  Finding direct evidence that the M star pair may be associated with SGR 1900+14 has proven elusive as summarized by \\citet{gue20}. Also difficult is a theoretical understanding of how isolated, albeit high mass, stars could play a role in the formation of a pulsating X--ray source, despite the presence of a high mass luminous blue variable (LBV) very near the SGR 1806--20 localization position, a remarkably similar situation to that for SGR 1900+14. Recent near-- and mid--infrared observations of SGR 1806--20 \\citep[F99]{fuc99}, however, have revealed the LBV to be only the most luminous member of a compact cluster of massive stars. Such proximate regions of recent star formation provide a natural location for the birth of such pulsating X--ray sources, which cannot be very old, without the need for invoking enormous space velocities from the nearest supernova remnants.  In this paper we present evidence for a similar compact cluster of high-mass stars which has heretofore been hidden in the glare of its brightest components, the pair of M5 supergiant stars. ",
        "conclusions": "If the cluster was the birthplace of SGR 1900+14, this essentially excludes SNR G42.8+0.6 as playing any role in the SGR. Although one can envision scenarios in which the SNR progenitor was ejected from the cluster by dynamical interaction or a much earlier supernova, this leaves the necessity of the neutron star having been kicked back to almost exactly its place of origin by the supernova that formed SNR G42.8+0.6 (since the cluster and SGR localizations are coincident), an unlikely coincidence both in space and timing. However, despite the association of G42.8+0.6 with SGR 1900+14 in the literature, there has been no  evidence supporting this association offered, such as the probablity of finding any SNR within a given distance, based on the number density of SNRs in the Galactic plane.  A more plausable scenario is one in which the cluster and associated dense gas/dust cloud hides a recent supernova. Evidence for this cloud comes from Figures 1 and 2 and the coincident extended strong far--infrared source indicating compact warm and extended cool dust (see V96). Optical extinction from this cloud combined with a 12--15~kpc distance explains why the supernova would not have been noticed historically. A very young SNR expanding into the dense wind--blown bubble due to mass loss from the supergiant stars in the cluster would be consistent with the otherwise unexplained persistent X--ray source at this position, RX J190717+0919.3 \\citep{hur96}. While no quiescent radio source is known at this position, a combination of self--absorption within the dense medium and rapid decay \\citep{rey88} could account for this. The supernova remnant evolutionary calculations of \\citet{tru99} indicate that for an ejecta mass of 1 M$_\\sun$, and an external density medium of 10 cm$^{-3}$, one finds characteristic sizes of $\\approx$ 1~pc at t~=~1000 yr; similar to that of the cluster dimensions at the estimated M supergiant distances.  The most likely position for the SGR itself is the \\citet{fra99} fading radio source located at $\\alpha$~=~19$^h$~07$^m$~14.33$^s$, $\\delta$~=~+09$^d$~19'~21.1\" (J2000), with positional accuracy of $\\pm$ 0.15 arcsec in each coordinate. With these astrometric positions we estimate the approximate distances from the center and edge of the cluster to the radio position as 12 arcsec (0.7--0.9~pc) and 5 arcsec (0.3--0.4~pc), respectively, based on the 12--15~kpc distance estimate. Thus, even at the extreme minimum age of the SGR based on the simplest magnetar physics ($\\approx$~700 yr; Kouveliotou et al. 1999) this implies a tangential velocity of $\\approx$ 420 km~s$^{-1}$ from the near edge of the cluster. While still an ample velocity for the runaway neutron star, it obviates the enormous space velocities implied by associating it with G42.8+0.6 \\citep{kou99}, which is about 12 arcmin away \\citep{hur99b}.  While an isolated instance of the compact, high mass cluster found at/near SGR 1900+14 would be dismissed as a chance superposition, its striking similarity to the cluster found near SGR 1806-20 by \\citet{fuc99} must be recognized. In that case, an LBV supergiant is found associated with a cluster of at least another four massive young stars enshrouded in a bright dust cloud as imaged by ISO and located only 7 arcsec from the SGR gamma--ray localization. With an approximate cluster radius of 8 arcsec and an estimated distance of 14.5 kpc, this implies a cluster radius of $\\approx$~0.6~pc. Now that similar compact clusters have been found near the positions of the two best studied SGRs (1806-20 and 1900+14) the possiblity that young SGR neutron stars have their origins in compact clusters should be considered seriously."
    },
    "0002/astro-ph0002060_arXiv.txt": {
        "abstract": "A historical perspective on the study of asymmetries in planetary nebulae (PNs) is presented. We also describe our ongoing work in high resolution spectroscopy of planetaries, and discuss some likely future directions for the study of asymmetrical PNs. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002256_arXiv.txt": {
        "abstract": "POINT-AGAPE is an Anglo-French collaboration which is employing the Isaac Newton Telescope (INT) to conduct a pixel-lensing survey towards M31. Pixel lensing is a technique which permits the detection of microlensing against unresolved stellar fields. The survey aims to constrain the stellar population in M31 and the distribution and nature of massive compact halo objects (MACHOs) in both M31 and the Galaxy.  In this paper we investigate what we can learn from pixel-lensing observables about the MACHO mass and fractional contribution in M31 and the Galaxy for the case of spherically-symmetric near-isothermal haloes. We employ detailed pixel-lensing simulations which include many of the factors which affect the observables, such as non-uniform sampling and signal-to-noise ratio degradation due to changing observing conditions.  For a maximum MACHO halo we predict an event rate in $V$ of up to 100 per season for M31 and 40 per season for the Galaxy. However, the Einstein radius crossing time is generally not measurable and the observed full-width half-maximum duration provides only a weak tracer of lens mass. Nonetheless, we find that the near-far asymmetry in the spatial distribution of M31 MACHOs provides significant information on their mass and density contribution. We present a likelihood estimator for measuring the fractional contribution and mass of both M31 and Galaxy MACHOs which permits an unbiased determination to be made of MACHO parameters, even from data-sets strongly contaminated by variable stars. If M31 does not have a significant population of MACHOs in the mass range $10^{-3}~\\sm - 1~\\sm$ strong limits will result from the first season of INT observations. Simulations based on currently favoured density and mass values indicate that, after three seasons, the M31 MACHO parameters should be constrained to within a factor four uncertainty in halo fraction and an order of magnitude uncertainty in mass ($90\\%$ confidence). Interesting constraints on Galaxy MACHOs may also be possible. For a campaign lasting ten years, comparable to the lifetime of current LMC surveys, reliable estimates of MACHO parameters in both galaxies should be possible. ",
        "introduction": "\\label{s1}  \\subsection{Conventional microlensing: landmarks and limitations}  The detection of the gravitational microlensing effect due to compact objects in the Galaxy is undoubtedly one of the great success stories in astrophysics over the past decade. Surveys have discovered around 20 candidates towards the Magellanic clouds and several hundred towards the Galactic Bulge \\cite{uda94,alc97,ala97,lass99,alc00}. Amongst these candidates a number of exotic lensing phenomena have been catalogued, such as parallax effects, binary lensing (including spectacular examples of caustic-crossing events), and finite source-size effects. These discoveries are facilitated by coordinated follow-up campaigns such as PLANET \\cite{alb98} and MPS \\cite{rhie99} which act on microlensing alerts broadcast by the survey teams. The absence of certain microlensing signals has also yielded a clearer insight into the nature of halo dark matter. The null detection of short duration events towards the Large Magellanic Cloud (LMC) by the EROS and MACHO surveys indicates that, for a range of plausible halo models, massive compact halo objects (MACHOs) within the mass interval $10^{-7} - 10^{-3}~\\sm$ provide less than a quarter of the dark matter \\cite{alc98}. This is an important result when set against the current insensitivity of other techniques to this mass range.  Despite these successes a number of unsolved problems remain. The optical depth measured towards the Galactic Bulge is at least a factor two larger than can be accommodated by theoretical models (e.g. Bissantz et al. 1997; Sevenster et al. 1999). Towards the LMC the rate of detected events is consistent with the discovery of a significant fraction of the halo dark matter. However, the implied lens mass range ($0.1 - 1~\\sm$) is not easily reconciled with existing constraints on baryonic dark matter candidates \\cite{carr94}, though the MACHOs need not necessarily be baryonic. Furthermore, the discovery of two possible binary caustic-crossing events towards the LMC and the Small Magellanic Cloud (SMC) has thrown into question the very existence of MACHOs.  Their caustic-crossing timescales, which provide an indicator of their line-of-sight position, seem to exclude either as being of halo origin, a statistically unlikely occurrence if the halo comprises a significant MACHO component \\cite{ker99}.  As a result, there is a growing body of opinion that all events observed so far towards the LMC and SMC may reside in the clouds themselves. However, this explanation is itself problematic because it requires that the clouds must either have a higher MACHO fraction than the Galaxy or comprise substantial but diffuse stellar components not in hydrodynamical equilibrium (Evans \\& Kerins 2000, and references therein).  These problems highlight two principal constraints on the ability of conventional microlensing experiments to determine the nature and distribution of MACHOs in the halo. The first limitation is their inefficiency in differentiating between lensing by MACHOs and self-lensing by the source population, since for most events one observes only a duration and a position on the sky. These observables are only weakly correlated with the location of the events along the line of sight. The second constraint is the limited number of suitable lines of sight through the halo. Conventional microlensing surveys require rich yet resolved stellar fields and are thus limited to just two lines of sight, the LMC and SMC, with which to probe MACHOs. The line of sight to the Galactic Bulge is dominated by bulge and disc lensing. The paucity of halo lines of sight, together with the rather weak dynamical and kinematical constraints on Galactic halo structure, also diminishes the prospect of being able to decouple information on the Galactic distribution function and MACHO mass function.  \\subsection{Beyond the Galaxy: a new target, a new technique}  The possibility of detecting MACHOs in an external galaxy, specifically M31, was initially explored by Crotts (1992) and by Baillon et al. (1993). Crotts (1992) pointed out that the high inclination of the disc of M31 would result in an asymmetry in the observed rate of microlensing if the disc is surrounded by a MACHO halo, as illustrated in Figure~\\ref{f1}. The fact that the M31 MACHO microlensing rate should be lower towards the near side of the disc than the far side, which lies behind a larger halo column density, means that the presence of MACHOs in M31 can be established unambiguously. In particular, neither variable stars nor stellar self-lensing events in the disc of M31 should exhibit near-far asymmetry. Additionally, the external vantage point serves to reduce systematic model uncertainties in two ways. Firstly, it permits a more accurate determination of the rotation curve and surface brightness profile than is possible for the Galaxy, which reduces the prior parameter space of viable galactic models. Secondly, it provides many independent lines of sight through the halo of M31, allowing the MACHO distribution across the face of the disc to be mapped and thus the halo distribution function to be constrained more or less directly.  \\begin{figure} \\begin{center} \\epsfig{file=fig01.ps,width=5cm,angle=270} \\end{center} \\caption{The principle of near-far asymmetry. The optical depth through the halo towards the far disc is larger than towards the near disc owing to the tilt of the disc confined within the spheroidal distribution of MACHOs. The distribution of Galaxy MACHOs, disc self-lensing events and variable stars does not exhibit asymmetry.} \\label{f1} \\end{figure}  As pointed out by Baillon et al. (1993), another appeal of directing observations towards more distant large galaxies like M31 is the increase in the number of potential source stars, more than a factor of one thousand over the number available in the LMC and SMC, and all confined to within a few square degrees.  However, this also presents a fundamental problem in that the source stars are resolved only whilst they are lensed (and even then only if the magnification is sufficiently large). The presence of many stars per detector pixel means it is often impossible to identify which is being lensed. Furthermore, the flux contribution of the unlensed stars dilutes the observed flux variation due to microlensing. Nonetheless, Baillon et al. (1993) determined from numerical simulations that the number of observable events, due to either the lensing of bright stars or high magnification events, is expected to be large. As a result of these studies, the Andromeda Galaxy Amplified Pixel Experiment (AGAPE) and another group, Columbia-VATT, commenced observing programs towards M31 \\cite{ans97,cro97}.  One of the biggest technical difficulties facing surveys which look for variable sources against unresolved stellar fields is how to distinguish between flux variations due to changing observing conditions and intrinsic variations due to microlensing or stellar variability. For example, changes in seeing induce variations in the detected flux within a pixel. One must also deal with the consequences of positional misalignment between exposures, spatial and temporal variations in the point spread function (PSF) and photometric variations due to atmospheric transparency and variable sky background.  AGAPE has employed the Pixel Method to cope with the changing observing conditions \\cite{ans97}.  AGAPE thoroughly tested this technique with a three-year campaign using the 2m Bernard Lyot telescope at Pic du Midi from 1994 to 1996 \\cite{ans97,ans99,ledu00}. Six fields covering about 100 arcmin$^2$ centred on the bulge of M31 were monitored. Whilst the field of view was insufficient to conclude much about the nature of MACHOs, 19 candidate events were detected, though it is still premature to rule out many of them being intrinsically variable sources, such as Miras or novae. One event, AGAPE~Z1, appears to be a convincing lensing candidate as its flux increase and colour are inconsistent with that of a Mira or nova \\cite{ans99}. A longer baseline is needed to determine how many of the other candidates are due to microlensing.  A major observing programme began on the 2.5m Isaac Newton Telescope (INT) in La~Palma in the Autumn of 1999, with a run of one hour per night for almost sixty nights over six months. The POINT-AGAPE collaboration is a joint venture between UK-based astronomers and AGAPE (where POINT is an acronym for ``Pixel-lensing Observations with INT'').  We are exploiting the 0.3~deg$^2$ field of view of the INT Wide-field Camera (WFC) to map the distribution of microlensing events across a large region of the M31 disc.  Our initial observations of M31 with the INT employed a $V$ filter and the simulations reported here have been undertaken with parameters appropriate to V-band observations. The strategy employed for the actual M31 monitoring campaign involves observations in three bands, $g$, $r$, and $i$ [very similar to the bands employed by SLOAN \\cite{fuk96}]. The multi-colour observations will improve our ability to discriminate against variable stars and the $gri$-filter plus CCD combination offers a significant improvement in sensitivity (the $g$-band zero-point is some 0.4 magnitudes fainter than that for $V$). The simulation parameters are thus somewhat conservative in this regard. The programme is being conducted in consort with the Microlensing Exploration of the Galaxy and Andromeda (MEGA) survey \\cite{cro99}, the successor program to Columbia-VATT. Whilst POINT-AGAPE and MEGA are sharing the data, different techniques are being employed to search for microlensing events. Henceforth we use the term {\\em pixel lensing}\\/ \\cite{gou96} to describe microlensing against unresolved stellar fields, regardless of the detection technique.  Whilst the technical viability of pixel lensing is now clearly established, a number of important theoretical issues are still outstanding. The principal concern is that the main observable in classical microlensing, the Einstein crossing time, is generally not accessible in pixel lensing. The Einstein crossing time is directly related to the lens mass, its transverse velocity and the observer--lens--source geometry. In pixel lensing the observed timescale depends upon additional factors, such as the local surface brightness and the source luminosity and magnification, so the dependence on lens parameters is much weaker than for classical microlensing.  The first detailed study of pixel lensing was undertaken by Gould (1996). He defined two regimes: a semi-classical regime in which the source star dominates the pixel flux and the observable timescale provides a fair tracer of the Einstein crossing time; and the ``spike'' regime where only high-magnification events are identified, and the timescales are only weakly correlated with the underlying Einstein crossing duration. Remarkably, Gould showed that, despite the loss of timescale information, in the spike regime one can still measure the microlensing optical depth. Using Gould's formalism, Han (1996) provided the first pixel event rate estimates for the M31 line of sight.  However, Gould's formalism assumes a fixed sampling rate and unchanging observing conditions. As such it is of limited applicability to a ground-based observing program. Gondolo (1999) has proposed an optical depth estimator based on the observed pixel event timescale. Whilst this estimator can be readily employed by a ground-based campaign, it is somewhat sensitive to the shape of the source luminosity function and is valid only to the extent that this can be taken to be the same for all source components.  More recently, Baltz \\& Silk (1999) derived expressions for the pixel rate and timescale distribution in terms of the observable timescale, rather than the Einstein crossing time. Again, their study assumes constant sampling and observing conditions, as would be the case for space-borne programmes.  Whilst these studies provide a solid foundation for predictions of pixel-lensing quantities (i.e. timescales, rates and optical depth), none of them address to what extent one can constrain galactic and lens parameters, in particular the MACHO mass, from pixel lens observables.  Gyuk \\& Crotts (2000) have shown that a reliable measure of the optical depth from pixel lensing can be used to probe the core radius and flattening of the M31 MACHO halo.  In this paper we quantitatively assess the degree to which the POINT-AGAPE campaign directed towards M31 will constrain the fractional contribution and mass of the MACHOs.  Since the answer inevitably depends upon the assumed galactic distribution function, we focus attention here on the simple case of spherically-symmetric near-isothermal halo models. The line of sight towards M31 is sensitive to two MACHO populations, our own and that in M31 itself, so we investigate the extent to which they can be distinguished and probed independently. We also model the expected background due to variable stars and lenses residing in the disc and bulge of M31.  The plan of the paper is as follows. In Section~\\ref{s2} we summarize the basic principles of pixel lensing, with emphasis on the differences between pixel lensing and classical microlensing. We describe our Monte-Carlo pixel-lensing simulations in Section~\\ref{s3}, including our event selection criteria and the incorporation of realistic sampling and observing conditions. In Section~\\ref{s4} we construct a reference model for the lens and source populations in the halo of the Galaxy and the halo, disc and bulge of M31, seeking consistency with the observed M31 rotation curve and surface brightness profiles. In Section~\\ref{s5} we present predictions for the POINT-AGAPE survey based on our simulations. In Section~\\ref{s6} we use the simulations to generate artificial data-sets and we investigate to what extent the MACHO mass and fractional contribution in the two galaxies can be recovered from the data. The results are summarized and discussed in Section~\\ref{s7}. ",
        "conclusions": "\\label{s7}  Pixel lensing is a relatively new and powerful method to allow microlensing searches to be extended to targets where the sources are unresolved. It heralds the possibility of detecting or constraining MACHO populations in external galaxies. Though pixel lensing is hampered by changes in observing conditions, which introduce spurious variations in detected pixel flux, techniques have been developed which minimize these variations to a level where genuine microlensing signals can be detected.  POINT-AGAPE and another team (MEGA) have embarked on a major joint observing programme using the Isaac Newton Telescope (INT) to monitor unresolved stars in M31 for evidence of pixel lensing due to MACHOs either in the Galaxy or M31 itself. Two techniques, the Pixel Method and difference imaging, are available to minimize flux variations induced by the changing observing conditions. In this paper we have assessed the extent to which the Pixel Method allows us to determine the mass and fractional contribution of MACHOs in both M31 and the Galaxy from pixel-lensing observables.  Our assessment takes account of realistic variations in observing conditions, due to changes in seeing and sky background, together with irregular sampling.  Pixel lensing observables differ from those in classical microlensing, where one targets resolved stellar fields, in that one is generally unable to measure the Einstein radius crossing time, $\\te$, of an event. The fact that the source stars are resolved only whilst they are lensed means that one is unable to determine their baseline luminosity, so neither the magnification nor the total duration of the event can be measured. As an alternative to $\\te$ one may measure the full-width half-maximum timescale, $\\tfw$, directly from the light-curve. However, this provides only a lower limit to the underlying event duration. Fortunately, M31 provides a signature which permits an unambiguous determination of whether or not MACHOs reside in its halo: near-far asymmetry \\cite{cro92}. If M31 is embedded in a dark spheroidal halo of MACHOs the high disc inclination should provide a measurable gradient in the observed pixel lensing rate. The strength of this signature depends both on the mass and fractional contribution of MACHOs in M31, as well as the level of ``contamination'' by variable stars, M31 stellar lensing events and foreground Galaxy MACHOs.  We have employed detailed Monte-Carlo simulations to estimate the timescale and spatial distributions of MACHOs in both our Galaxy and M31 for spherically-symmetric near-isothermal halo models. We also model the lensing contribution due to disc and bulge self-lensing. The expected number of M31 MACHOs for our two INT fields peaks at about 100 events for $\\sim 0.01~\\sm$ MACHOs, the Galaxy MACHO contribution being about half as large. For a given mass and halo fraction we expect to detect about an order of magnitude more events than current conventional surveys targeting the LMC.  The timescale distributions for Galaxy and M31 MACHOs are practically identical because of the symmetry of the microlensing geometry. Our simulations also confirm that $\\tfw$ is less strongly correlated with lens mass than $\\te$. For our sampling we find that, empirically, $\\langle \\tfw \\rangle \\propto \\langle \\te \\rangle^{1/2} \\propto m^{1/4}$ for lens mass $m$.  Sampling introduces a significant bias in the duration of detected events with respect to the underlying average for very massive and very light MACHOs.  Our simulations clearly show the near-far asymmetry in the M31 MACHO spatial distribution. However, the presence of the foreground Galaxy MACHOs makes its measurement more difficult. We also find that the distribution of very massive MACHOs is noticeably more centrally concentrated than that of less massive lenses.  Stellar self-lensing events are found to be mostly confined to within the inner 5~arcmin of the M31 disc, and are mostly due to bulge self-lensing. Their tight spatial concentration means that they do not pose a serious contamination problem for analysis of the Galaxy and M31 MACHO populations.  We have constructed a maximum likelihood estimator which uses timescale and position observables to simultaneously constrain the MACHO mass and halo fraction of both M31 and the Galaxy. The statistic is devised to be robust to data-set contamination by variable stars. We find that M31 MACHO parameters can be reliably constrained by pixel lensing. For simulated INT data-sets we find pixel-lensing constraints on the M31 halo to be comparable to those obtained for the Galaxy halo by the conventional microlensing surveys. Even with severe contamination from variable stars the M31 MACHO parameters are well determined within three years. In particular, if there are few MACHOs in M31 this should become apparent after just one season of data collection, even if as many as a hundred variable stars pass the microlensing selection criteria, because of the absence of near-far asymmetry. Pixel lensing is less sensitive to Galaxy MACHO parameters. Our simulations indicate that we require at least three times as much observing time in order to produce comparable constraints on Galaxy MACHO parameters. If the spatial distribution of variable stars closely follows that of Galaxy MACHOs, then it may become very difficult to reliably constrain Galaxy MACHO parameters.  The work presented here clearly demonstrates that a vigorous monitoring campaign on a 2m class telescope with a wide-field camera can identify and characterize MACHOs in M31. We now have the opportunity to unambiguously establish the existence or absence of MACHOs in an external galaxy.  The advantage of targeting M31 over our own Galaxy is that we have many lines of sight through the halo of M31 and a clear signature with which to distinguish M31 MACHOs from stellar self-lensing, the primary source of systematic uncertainty for Galaxy halo microlensing surveys.  M31 therefore represents one of the most promising lines of sight for MACHO studies."
    },
    "0002/astro-ph0002126_arXiv.txt": {
        "abstract": "Recently it has been proposed that the main contributor to the dark energy of the Universe is a dynamical, slow evolving, spatially inhomogeneous scalar field, called the quintessence. We investigate the behavior of this scalar field at galactic level, trying it as the dark matter in the halos of galaxies. Using an exact solution of the Einstein's equations, we find an excellent concordance between our results and observations.\\\\ ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002310_arXiv.txt": {
        "abstract": "We use large-scale cosmological simulations to estimate the mass-to-light ratio of galaxy systems as a function of scale, and compare the results with observations of galaxies, groups, clusters, and superclusters of galaxies. We find remarkably good agreement between observations and simulations. Specifically, we find that the simulated mass-to-light ratio increases with scale on small scales and flattens to a constant value on large scales, as suggested by observations. We find that while mass typically follows light on large scales, high overdensity regions --- such as rich clusters and superclusters of galaxies --- exhibit higher $M/L_{\\rm B}$ values than average, while low density regions exhibit lower $M/L_{\\rm B}$ values; high density regions are thus \\emph{antibiased} in $M/L_{\\rm B}$, with mass more strongly concentrated than blue light. This is true despite the fact that the galaxy mass density is unbiased or positively biased relative to the total mass density in these regions. The $M/L_{\\rm B}$ antibias is likely due to the relatively old age of the high density regions, where light has declined significantly since their early formation time, especially in the blue band which traces recent star formation. Comparing the simulated results with observations, we place a powerful constraint on the mass density of the universe; using, for the first time, the entire observed mass-to-light function, from galaxies to superclusters, we find $\\Omega =0.16\\pm0.05$. ",
        "introduction": "One of the oldest - and simplest - techniques for estimating the mass density of the universe is the mass-to-light method. In this method, the average ratio of the observed mass to light of the largest possible systems is used; assuming it is a fair sample, it can then be multiplied by the total luminosity density of the universe to yield the universal mass density. When the method is applied to rich clusters of galaxies --- the largest virilized systems for which a mass has been reliably determined --- the total mass density of the universe adds up to only $\\Omega \\simeq 0.2\\,\\ $% (where $\\Omega $ is the mass density in units of the critical density) (Zwicky 1957, Abell 1965, Ostriker, Peebles \\& Yahil 1974, Bahcall 1977, Faber \\& Gallagher 1979, Trimble 1987, Peebles 1993, Bahcall, Lubin \\& Dorman 1995, Carlberg \\emph{et al.} 1996, 1997, and references therein). A fundamental assumption in this determination, however, is that the mass-to-light ratio ($M/L$) of clusters is a fair representation of the universal value. If the mass-to-light ratio of clusters is larger or smaller than the universal mean, then the resulting $% \\Omega $ will be an over- or under- estimate, respectively. It is not clear whether this classic assumption of an unbiased representation by clusters is correct. More generally, if mass follows light (i.e., galaxies) on large scales --- thus $M/L\\simeq $ constant --- the galaxy distribution is considered to be unbiased with respect to mass; if mass is distributed more broadly than light, as is generally believed, then the galaxy distribution is biased (i.e., more clustered) with respect to mass, and the above determination of $\\Omega $ is an underestimate. \\ We investigate these questions of cluster representation and bias, and the impact they have on the measurement of $\\Omega $.  Observations of galaxies, groups and clusters of galaxies suggest that $M/L$ increases as a function of scale up to scales of hundreds of kiloparsecs (Schwarzschild 1954, Rubin \\& Ford 1970, Roberts \\& Rots 1973, Ostriker \\emph{et al.} 1974, Einasto \\emph{et al.} 1974, Davis \\emph{et al.} 1980, Trimble 1987, Gramann 1990, Zaritzky \\emph{et al.} 1993, Fischer \\emph{et al }$1999$), but then flattens out and remains approximately constant on larger scales (Bahcall, Lubin \\& Dorman 1995). In the modern context we normally interpret this fact as indicating that luminous galaxies are more concentrated in peak density regions than the dark matter because baryons are dissipational. The shape and amplitude of the mass-to-light function --- that is, the dependence of $M/L$ on scale, $(M/L)(R)$ --- can place powerful constraints on the amount and distribution of dark matter in the universe, as well as on the amount of bias and its dependence on scale. The $M/L$ function thus provides a direct, model-independent census of the total mass density of the universe.  What is the expected dependence of $M/L$ on scale? In this paper we investigate this question  using large-scale, high resolution hydrodynamic cosmological simulations that contain dark matter and gas, and compare the results with observations. We find an excellent agreement between models and observations in the shape of the $M/L$ function; both data and models show an increase on small scales (hundreds of kpc) and a flattened $(M/L)(R)$ distribution on large scales. We use the comparison between data and simulations to determine the mass density of the universe. The amount of bias and its dependence on scale are also revealed. \\ We find that clusters of galaxies are mildly \\emph{antibiased, } in the sense that mass is more concentrated than light on average. Previous determinations of $\\Omega$ using clusters of galaxies have thus \\emph{overestimated} $\\Omega $ due to this unaccounted antibias. The present investigation attempts to provide an unbiased determination of $\\Omega$ using, for the first time, the entire observed mass-to-light function. The above results do not disagree with previous estimates that the mass density of galaxies is unbiased or positively biased with respect to the total mass density in the high density regions; it is the light density that is shown here to be antibiased. ",
        "conclusions": "\\qquad We use large-scale cosmological simulations to determine the expected mass-to-light ratio of galaxy systems and its dependence on scale. The $(M/L_{\\rm B})(R)$ function is investigated from small scales of galaxies ($R \\simeq 20$ kpc) to large scales ($R \\simeq 60h^{-1}$Mpc), and compared with observations of galaxies, groups, clusters, and superclusters. We use the results to evaluate the amount of bias on different scales (i.e., how mass traces light), and use the comparison with observations to determine the mass density of the universe, $\\Omega $.  We find the following results:  \\begin{enumerate} \\item In high density regions the galaxy blue light is antibiased (i.e., lower) relative to the total mass density (while the galaxy mass density is not). This is due to the old age of the high density systems which leads to a relative decrease in their present-day luminosity, especially in the blue band that traces recent star formation.  \\item The shape of the simulated $(M/L_{\\rm B})(R)$ function is in excellent agreement with observations. The simulated $M/L_{\\rm B}(R)$ function increases with scale on small scales and flattens on large scales, where $M/L_{\\rm B}$ reaches a constant value, as observed. The mean flattening of $(M/L_{\\rm B})(R)$ on large scales indicates that, on average, mass follows light on large scales (i.e., $M \\propto L$).  \\item Even though $M/L_{\\rm B}$ is approximately constant on large scales, we find that the actual value of $M/L_{\\rm B}$ depends on the local mass overdensity, $\\Delta \\rho /\\rho (<R)$, at a given scale. \\ High overdensity regions exhibit higher $M/L_{\\rm B}$ ratios than lower density regions. \\ The difference can typically be a factor of 2 to 3, consistent with observations of groups and clusters of galaxies (representing low and high density regions, respectively). The dependence of $M/L_{\\rm B}(R)$ on overdensity indicates that high density regions such as rich clusters and superclusters are relatively \\emph{antibiased% } - they exhibit higher than average $M/L_{\\rm B}$ values, implying that mass is more concentrated than light in these regions (see 1 above). In the blue luminosity band, the cluster $M/L_{\\rm B}$ antibias is $b_{\\rm cl}^{^{M/L_{\\rm B}}}=<M/L_{\\rm B}>_o/<M/L_{\\rm B}>_{\\rm cl}=0.75\\pm0.15$.  \\item We find that the $(M/L_{\\rm B})(R)$ function of high density regions is traced well by $(M/L_{\\rm B})(R)$ of old (elliptical) galaxies; low density regions are traced well by young (spiral) galaxies. \\ These results are consistent with observations.  \\item We determine the mass density of the universe by fitting the simulated $(M/L_{\\rm B})(R)$ function to observations. The best fit $\\Omega$ is lower than previous estimates based on cluster $M/L$ values because of the antibias discussed above as well as the more robust use of the entire $M/L$ function --- not just clusters --- in constraining $\\Omega$. We find a best-fit value of $\\Omega=0.16\\pm0.05$ (with an additional estimated uncertainty of $\\pm0.03$ for possible additional systematics); this value provides a remarkably good match to the data for galaxies, groups, clusters, and superclusters. \\ The results are independent of the details of the models and provide a powerful measure of $\\Omega$. The only significant uncertainty we are aware of is due to the possibility that current observations may systematically underestimate the global mean luminosity density of the universe. This will produce a corresponding underestimate in our computation of $\\Omega$ unless there was also a corresponding underestimate in the luminosity of groups, clusters, and superclusters of galaxies. \\end{enumerate}"
    },
    "0002/astro-ph0002036_arXiv.txt": {
        "abstract": "The rate of mass accumulation due to galaxy merging depends on the mass, density, and velocity distribution of galaxies in the near neighborhood of a host galaxy. The fractional luminosity in kinematic pairs combines all of these effects in a single estimator which is relatively insensitive to population evolution. Here we use a k-corrected and evolution compensated volume-limited sample having an R-band absolute magnitude of $M_R^{k,e} \\le -19.8+5\\log{h}$ mag drawing about 300 redshifts from CFGRS and 3000 from CNOC2 to measure the rate and redshift evolution of merging.  The combined sample has an approximately constant co-moving number and luminosity density from redshift 0.1 to 1.1 ($\\Omega_M=0.2, \\Omega_\\Lambda=0.8$); hence, any merger evolution will be dominated by correlation and velocity evolution, not density evolution. We identify kinematic pairs with projected separations less than either 50 or 100 \\hkpc\\ and rest-frame velocity differences of less than 1000\\kms.  The fractional luminosity in pairs is modeled as $f_L(\\Delta v,r_p,M_r^{ke})(1+z)^{m_L}$ where $[f_L,m_L]$ are $[0.14\\pm0.07,0\\pm1.4]$ and $[0.37\\pm0.7,0.1\\pm0.5]$ for $r_p\\le 50$ and 100\\hkpc, respectively ($\\Omega_M=0.2, \\Omega_\\Lambda=0.8$). The value of $m_L$ is about 0.6 larger if $\\Lambda=0$.  To convert these redshift space statistics to a merger rate we use the data to derive a conversion factor to physical space pair density, a merger probability and a mean in-spiral time. The resulting mass accretion rate per galaxy ($M_1,M_2\\ge 0.2 M_\\ast$) is $0.02\\pm0.01(1+z)^{0.1\\pm0.5} M_\\ast~{\\rm Gyr}^{-1}$.  Present day high-luminosity galaxies therefore have accreted approximately $0.15M_\\ast$ of their mass over the approximately 7 Gyr to redshift one.  Since merging is likely only weakly dependent on host mass, the fractional effect, $\\delta M/M \\simeq 0.15M_\\ast/M$, is dramatic for lower mass galaxies but is, on the average, effectively perturbative for galaxies above $M_\\ast$. ",
        "introduction": "Merging is a fundamental mode of stellar mass addition to galaxies. Moreover, merging brings in new gas and creates gravitational disturbances that enhance star formation or fuel a nuclear black hole. The general process of substructure infall may be the rate fixing process for the buildup of a galaxy's stars and consequently may largely regulate its luminosity history.  Gravitational forces on relatively large scales dominate merger dynamics which allows direct observation of the mechanism, although with the considerable complication that dark matter dominates the mass.  N-body simulations \\citep{tt,bh} give the detailed orbital evolution, morphological disturbances and eventual outcomes of the encounters of pairs of galaxies.  The purpose of this paper is to estimate the rate of mass gain per galaxy due to mergers over the redshift zero to one interval.  Our primary statistic is the fractional luminosity in close kinematic pairs, which is readily related to n-body simulations and sidesteps morphological interpretation. This approach provides a clear sample definition which is closely connected to the large scale dynamics of merging. In common with all merger estimates it requires an estimate of the fraction of the pairs that will merge and a mean time to merger.  The number of kinematic pairs is proportional to the volume integral at small scales of the product of two-point correlation function, $\\xi$, and the luminosity function (LF).  The high luminosity galaxies appear to be evolving purely in luminosity \\citep{cfrs_lf,huan_lf}, which can be easily compensated. The measured evolution of $\\xi$ suggests that the density of physical pairs should not vary much with redshift, $(1+z)^{0\\pm1}$ \\citep{cfrsxi,kkeck,cnoc_xi}. This inference is in notable contrast with the pair counts or morphological typing approaches to merger estimation \\citep{zk,cpi,ye_pairs,patton_cnoc1,cfrs_mg}, which suggest that merging rate by number varies as $(1+z)^{3\\pm1}$. HST photometric pairs, with no redshift information leads to a dependence of  $(1+z)^{1.2\\pm0.4}$ \\citep{mdss}.  In the next section we combine the Caltech Faint Galaxy Redshift Survey (CFGRS) and the Canadian Network for Observational Cosmology field galaxy survey (CNOC2) from which we construct evolution compensated, volume-limited, subsamples. In Section 3 we measure the fractional luminosity in 50 and 100\\hkpc\\ companions as a function of redshift. The CNOC2 sample is used in Section 4 to relate this wide pair sample to a close pair sample which is more securely converted into a mass merger rate. Section 5 discusses our conclusions. We use $H_0= 100h\\kmsm$, $\\Omega_M=0.2$ in open and flat cosmologies. ",
        "conclusions": "Our main observational result is that for galaxies with $M_R^{k,e}\\le -19.8+5\\log{h}$ mag, the fraction of galaxy luminosity in 50\\hkpc\\ wide kinematic pairs is about 14\\%, with no noticeable redshift dependence over the redshift zero to one range. This implies an integrated mass accretion rate of about 2\\% of $L_\\ast$ per Gyr per galaxy for merging galaxies having $L\\ge 0.2L_\\ast$.  Our rate is uncertain at about the factor of two level due to uncertainty in the dynamical details of merging for our sample definitions.  This merger rate implies a 15\\% mass increase in an $M_\\ast$ galaxy since redshift one. If the correlations of lower luminosity galaxies are only somewhat weaker than these \\citep{roysoc} then the same $0.15M_\\ast$ merged-in mass causes a 50\\% mass increase in a 0.3$M_\\ast$ galaxy.  There are several issues that require further investigation. First, the rate of merging of similarly selected kinematic pairs should be studied in appropriately matched n-body experiments to better determine the orbital timescales. Second, the absence of a redshift dependence of $\\sigma_{12}$ and $v_{mg}$ needs to be observationally checked. Third, the connection between close kinematic pairs and morphologically disturbed galaxies, which does conform to the kinematic pair predictions at low redshift \\citep{patton_ssrs2}, needs to be better understood at high redshift."
    },
    "0002/astro-ph0002200_arXiv.txt": {
        "abstract": "We present new techniques to produce IRAS 12 $\\mu$m samples of galaxies and stars. We show that previous IRAS 12 $\\mu$m samples are incompatible for detailed comparison with ISO surveys and review their problems. We provide a stellar infrared diagnostic diagram to distinguish galaxies from stars without using longer wavelength IRAS colour criteria and produce complete 12 $\\mu$m samples of galaxies and stars. This new technique allows us to estimate the contribution of non-dusty galaxies to the IRAS 12 $\\mu$m counts and produce a true local mid-infrared extragalactic sample compatible with ISO surveys. We present our initial analysis and results. ",
        "introduction": "The recent ISO mission has produced a number of deep mid-infrared extragalactic surveys [1,2,3,4] many of which are presented elsewhere in these proceedings. In order to accurately evaluate the apparent source evolution found in these surveys it is essential to have a stable and exact local infrared picture that is compatible with ISO surveys. ",
        "conclusions": ""
    },
    "0002/astro-ph0002493_arXiv.txt": {
        "abstract": "We present near-infrared K-band spectroscopy of 21 elliptical or cD Brightest Cluster Galaxies (BCGs), for which we have measured the strength of the 2.293~$\\mu$m CO stellar absorption feature.  We find that the strength of this feature is remarkably uniform among these galaxies, with a smaller scatter in equivalent width than for the normal elliptical population in the field or clusters. The scatter for BCGs is 0.156~nm, compared with 0.240~nm for Coma cluster ellipticals, 0.337~nm for ellipticals from a variety of other clusters, and 0.422~nm for field ellipticals. We interpret this homogeneity as being due to a greater age, or more uniform history, of star formation in BCGs than in other ellipticals; only a small fraction of the scatter can be due to metallicity variations, even in the BCGs. Notwithstanding the small scatter, correlations are found between CO strength and various galaxy properties, including R-band absolute magnitude, which could improve the precision of these galaxies as distance indicators in measurements of cosmological parameters and velocity flows. ",
        "introduction": "The equivalent widths of CO absorption features for the sample of BCGs observed in this study are presented in Table 1.  The data included in this table are Abell \\shortcite{ab:58} catalogue numbers (column 1), BCG names (column 2), CO$_{EW}$ values with 1--$\\sigma$ errors (column 3), recession velocity in kms$^{-1}$ (column 4), absolute R-band magnitude corresponding to the metric luminosity $L_m$ (column 5), structure parameter ($\\alpha$) (column 6) and the magnitude residual relative to the best-fit $L_m$--$\\alpha$ relation (column 7), (columns 4--7 are all taken from Lauer \\& Postman \\shortcite{la:94}, who assumed a Hubble constant of 80~kms$^{-1}$Mpc$^{-1}$).  Columns 8 and 9 contain velocity dispersions and Mg$_2$ metallicity indices, where available, from Faber et al. \\shortcite{fa:89}.  \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Photometric and spectroscopic parameters for 21 brightest cluster galaxies} \\begin{tabular}{llllcclcl}\\\\ & & & & & & & &\\\\ Abell$\\#$   & Galaxy name & CO$_{EW}$(nm)& V$_{rec}$ & M$_R$ & $\\alpha$ & dM$_{\\alpha}$ & $\\sigma$ (kms$^{-1}$) & Mg$_2$ \\\\  & & & & & & &\\\\  496  & MCG-02-12-039 &   3.36$\\pm$0.20 &  9893 & -22.579 &  0.786 & 0.103 &  274 & --\\\\ 533  &  --           &   3.55$\\pm$0.19 & 14365 & -22.397 &  0.500 & 0.012 &  --  & --\\\\ 539  & MCG+01-14-019 &   3.36$\\pm$0.18 &  9682 & -22.366 &  0.505 & 0.050 &  309 & --\\\\ 548  & ESO488-G033   &   3.23$\\pm$0.18 & 11848 & -22.490 &  0.493 &-0.093 &  --  & --\\\\ 569  & NGC2329       &   3.15$\\pm$0.18 &  5724 & -22.357 &  0.472 & 0.003 &  271 & 0.269\\\\ 576  &  --           &   3.42$\\pm$0.18 & 12072 & -22.051 &  0.296 &-0.083 &  --  & --\\\\ 634  & UGC4289       &   3.20$\\pm$0.18 &  8135 & -22.248 &  0.497 & 0.156 &  241 & --\\\\ 671  & IC2378        &   3.44$\\pm$0.19 & 14970 & -22.961 &  0.713 &-0.310 &  313 & --\\\\ 779  & NGC2832       &   3.48$\\pm$0.18 &  6867 & -23.011 &  0.596 &-0.465 &  354 & 0.340\\\\ 912  &  --           &   3.23$\\pm$0.19 & 13572 & -21.984 &  0.423 & 0.283 &  --  & --\\\\ 957  & UGC5515       &   3.39$\\pm$0.19 & 13438 & -22.882 &  0.760 &-0.207 &  350 & --\\\\ 999  & MCG+02-27-004 &   3.34$\\pm$0.18 &  9749 & -22.328 &  0.443 &-0.021 &  --  & --\\\\ 1016 & IC613         &   3.28$\\pm$0.18 &  9705 & -22.112 &  0.436 & 0.181 &  --  & --\\\\ 1060 & NGC3311       &   3.39$\\pm$0.24 &  3704 & -22.392 &  0.845 & 0.297 &  192 & 0.297\\\\ 1142 & IC664         &   3.34$\\pm$0.18 & 10118 & -22.367 &  0.544 & 0.110 &  --  & --\\\\ 1656 & NGC4889       &   3.55$\\pm$0.18 &  6497 & -23.106 &  0.612 &-0.541 &  404 & 0.359\\\\ 2147 & UGC10143      &   3.01$\\pm$0.18 & 10384 & -22.374 &  0.666 & 0.244 &  303 & --\\\\ 2162 & NGC6086       &   3.25$\\pm$0.18 &  9547 & -22.594 &  0.503 &-0.179 &  325 & 0.344\\\\ 2197 & NGC6173       &   3.62$\\pm$0.25 &  8800 & -22.988 &  0.592 &-0.466 &  295 & 0.332\\\\ 2199 & NGC6166       &   3.23$\\pm$0.25 &  9348 & -22.748 &  0.777 &-0.067 &  320 & 0.340\\\\ 2634 & NGC7720       &   3.60$\\pm$0.25 &  9141 & -22.662 &  0.643 &-0.065 &  305 & 0.339\\\\  \\end{tabular} \\end{minipage} \\end{table*}    \\begin{figure} \\centerline{\\psfig{figure=bcg_fig2a.eps,width=0.5\\textwidth,angle=270}} \\centerline{\\psfig{figure=bcg_fig2b.eps,width=0.5\\textwidth,angle=270}} \\centerline{\\psfig{figure=bcg_fig2c.eps,width=0.5\\textwidth,angle=270}} \\caption{Histograms of CO$_{EW}$ distributions BCGs (hashed regions), overlaid on field and group galaxies (a), cluster galaxies (b), and Coma cluster galaxies (c).} \\end{figure}  We find the mean CO$_{EW}$ value for the 21 BCGs (3.35$\\pm$0.03) to be effectively identical to that of the Coma cluster ellipticals (3.37$\\pm$0.04) \\cite{mo:99}, and to that of the 31 ellipticals from a range of clusters discussed by James \\& Mobasher \\shortcite{ja:99} (3.29$\\pm$0.06).  The cluster and BCG distributions lie between the distributions of `isolated' and `group' field ellipticals discussed by James \\& Mobasher \\shortcite{ja:99} and shown in Fig. 2a.  The major difference between the BCG CO$_{EW}$ values and those of other ellipticals is the remarkably small range in the former: the standard deviation for BCGs is 0.156, compared to 0.240 for Coma ellipticals, 0.337 for general cluster galaxies, and 0.422 for cluster plus field ellipticals.  Indeed, the scatter in BCG CO absorption strengths is that predicted from the error estimates on the individual CO$_{EW}$ values, and so the intrinsic scatter may be much smaller still.  Given the small number of BCG galaxies, a Kolmogorov-Smirnov test cannot distinguish between the distributions of BCG and cluster or Coma galaxies in Figs. 2b and 2c, but there is less than 10\\% chance that the BCGs are drawn from the same parent population as all the non-BCG ellipticals, and less than 1\\% chance that they are from the same population as field ellipticals (Fig. 2a).  The BCGs are drawn from a much narrower region of the galaxy luminosity function than are the comparison samples in Fig. 2, which could affect the interpretation of this result.  The 21 BCGs have a range in M$_R$ of -22.0 to -23.1, little more than a magnitude. R-band photometry is not available for all the comparison galaxies, but good estimates can be made from published optical and near-IR photometry, leading to an estimated range of M$_R$ of -19.4 to -22.8 for the Coma cluster ellipticals, and -19.8 to -22.5 for the field and cluster ellipticals discussed by James and Mobasher \\shortcite{ja:99}.  We investigated whether the differences in CO$_{EW}$ scatter shown in Fig. 2 result from these differences in luminosity range by regressing CO$_{EW}$ on absolute magnitude, and studying the distributions of CO$_{EW}$ residuals about the best-fit lines.  The distributions of these residuals are shown in Fig. 3.  The dashed, diagonally shaded columns represent the residuals for the BCGs; the thick, dotted columns are those for the Coma cluster galaxies; and the solid lines represent the residuals for the field, group and cluster galaxies from James and Mobasher \\shortcite{ja:99}. The standard deviations of the CO$_{EW}$ residuals are 0.133~nm for the BCGs, 0.221~nm for the Coma ellipticals, and 0.419~nm for the cluster plus field ellipticals. This reinforces the conclusion from Fig. 2 that the BCGs have substantially more homogeneous CO strengths than the other elliptical galaxies studied, and this result does not appear to be a selection effect caused by the small luminosity range of the BCGs.  \\begin{figure} \\centerline{\\psfig{figure=bcg_fig3.eps,width=0.5\\textwidth,angle=270}} \\caption{Histograms of residuals of CO$_{EW}$ when regressed on absolute magnitude, for BCGs (hashed regions with dashed lines), overlaid on residuals for Coma ellipticals (thick dotted lines) and general field and cluster ellipticals (solid lines).} \\end{figure}   This uniformity in CO$_{EW}$ values is the main result of this paper, and it is important to consider what it implies in terms of differences between BCGs and other ellipticals.  Both high metallicity and recency of star formation are expected to increase CO$_{EW}$ values. The effect of metallicity on CO$_{EW}$ values can be estimated for the galaxies with measured Mg$_2$ indices, using the following method.  From Fig. 37 of Worthey \\shortcite{wo:94}, a change in [Fe/H] from -0.25 to 0.00 changes the Mg$_2$ index from 0.216 to 0.258, and the change is approximately linear over the modelled range.  Thus we infer a relation of the form $$\\delta Mg_2 = 0.168~\\delta [Fe/H].$$ Doyon et al. \\shortcite{do:94} find a relation between [Fe/H] and their CO index CO$_{sp}$, $$\\delta CO_{sp} = 0.11~\\delta [Fe/H],$$ and from the definitions in Puxley et al. \\shortcite{pu:97} it is straightforward to convert from the index CO$_{sp}$ to CO$_{EW}$. Then, the measured scatter in Mg$_2$ index of 0.029 for the BCGs should cause a scatter of 0.060~nm in CO$_{EW}$, 38\\% of the observed scatter.  For Coma ellipticals, the measured Mg$_2$ scatter is 0.024, equivalent to a scatter of 0.049~nm in CO$_{EW}$, 20\\% of that observed, and for the field and cluster sample, the Mg$_2$ scatter is 0.030, and the predicted CO$_{EW}$ scatter 0.062~nm, 15\\% of that observed. Note also that the scatters in Mg$_2$ values are very similar in the three subsamples, whereas they have very different CO$_{EW}$ distributions. Thus, we conclude that metallicity differences have little effect on the measured CO$_{EW}$ values for the elliptical galaxies studied here, and propose that star formation history is the dominant factor causing the larger scatter for non-BCG ellipticals. If so, the differences in the distributions of CO$_{EW}$, shown in Fig. 2 would be the result of wider variations in star formation history for general field and cluster ellipticals than for the BCGs. This indicates that BCGs formed their stars very early; if there has been more recent star formation in these galaxies then the rate of star formation as a function of epoch must have been very uniform from galaxy to galaxy.  Given the narrow range in BCG CO$_{EW}$ values, it is unrealistic to expect very strong correlations with other BCG parameters. Nevertheless, Fig. 4 does show a good correlation with absolute R-band magnitude in a 10~kpc metric aperture, M$_R$, with a correlation coefficient of 0.51 and a probability of 98.4\\% that this represents a true correlation (i.e. 1.6\\% probability that it could arise by chance). This is significant enough to be useful as a distance indicator: the scatter in M$_R$ for the 21 galaxies observed is 0.326~mag, which reduces to 0.280~mag when the M$_R$ values are corrected for the CO$_{EW}$ effect. The slope of the regression line of M$_R$ on CO$_{EW}$ is somewhat smaller than that for the trend in M$_K$ (total K-band absolute magnitude) vs CO$_{EW}$ for Coma cluster ellipticals \\cite{mo:99}, at -1.1$\\pm$0.5~mag/nm c.f. -1.6$\\pm$0.7~mag/nm for the Coma galaxies.  However, this difference is not statistically significant ($\\sim$0.6$\\sigma$).  It is not possible to determine whether the absolute magnitude--CO$_{EW}$ relations are consistent for the various samples because of the lack of homogeneous photometry, and the consequent need for large and uncertain colour and aperture corrections.   \\begin{figure} \\centerline{\\psfig{figure=bcg_fig4.eps,width=0.5\\textwidth,angle=270}} \\caption{Absolute R-band magnitude versus CO Equivalent Width for the 21 Brightest Cluster Galaxies. } \\end{figure}    Similarly, there is a strong correlation between CO$_{EW}$ and residuals (dM$_{\\alpha}$) about the relation of M$_R$ with structure parameter $\\alpha$ \\cite{ho:80} (Fig. 5), in the sense that galaxies with high CO$_{EW}$ tend to be bright relative to the mean relation (correlation coefficient 0.60, significance 99.6\\%). The residuals (dM$_{\\alpha}$) are reduced from 0.243~mag to 0.195~mag by correcting for the trend with CO$_{EW}$ shown in Fig. 5.  There is no correlation between CO$_{EW}$ and the structure parameter $\\alpha$ itself.  \\begin{figure} \\centerline{\\psfig{figure=bcg_fig5.eps,width=0.5\\textwidth,angle=270}} \\caption{Structure parameter residuals versus CO Equivalent Width for the 21 Brightest Cluster Galaxies. } \\end{figure}  Given the trends found in Figs. 4 \\& 5, it is instructive to explore if these effects could cause the putative streaming flow signal detected by Lauer \\& Postman \\shortcite{la:94} using the full sample of BCGs. However, we find no significant correlation between CO$_{EW}$ and direction on the sky (Fig. 6).  This implies that the Lauer \\& Postman \\shortcite{la:94} apparent detection of a bulk flow was not an artefact of differing stellar populations between sample galaxies, although we have of course only looked at a small fraction (18\\%) of their sample.  \\begin{figure} \\centerline{\\psfig{figure=bcg_fig6.eps,width=0.5\\textwidth,angle=270}} \\caption{CO$_{EW}$ versus the cosine of the angle between galaxy direction and the apex of the Lauer \\& Postman\\shortcite{la:94} streaming motion, for the 21 Brightest Cluster Galaxies. } \\end{figure}   The trends in figures 4 \\& 5 are both in the sense that brighter galaxies have higher indices: that in Fig. 4 could be a consequence of a metallicity--absolute magnitude relation, and there is indeed evidence of a weak correlation of CO$_{EW}$ with metallicity (Fig. 7). The correlation coefficient here is 0.52, but the relation has only 80\\% significance due to only 8 BCGs having tabulated Mg$_2$ values. Fig. 7 also shows the trend in CO$_{EW}$ with metallicity for 31 Coma cluster galaxies \\cite{mo:99}. The BCGs clearly lie at higher mean metallicity than do the Coma cluster ellipticals (mean Mg$_2$ values 0.328 for the BCGs and 0.294 for the Coma ellipticals).  Using the relation between Mg$_2$ and [Fe/H] from Worthey \\shortcite{wo:94}, and that between [Fe/H] and CO absorption strength found by Doyon et al. \\shortcite{do:94}, this predicts a difference in mean CO$_{EW}$ of 0.05~nm, compared to the observed difference of 0.04$\\pm$0.07~nm in the same sense. The small variation found here confirms the weakness of the dependence of CO$_{EW}$ on metallicity, as found by Doyon et al. \\shortcite{do:94}, at least at the high metallicity values typical of centres of bright galaxies, and also confirms our earlier conclusion that star formation history is the dominant effect in determining CO$_{EW}$ strength. This is further confirmed by the wide spread in the CO$_{EW}$ values of isolated and group ellipticals \\cite{ja:99}, with no corresponding change in their metallicity.   Finally, we find a weak correlation of CO$_{EW}$ with velocity dispersion for 14 galaxies with data in Table 1 (Fig. 8). This corresponds to a correlation coefficient of 0.294, and is significant at the 69\\% level. The slope and correlation coefficient are the same as was found for 31 Coma cluster ellipticals, also plotted in Fig. 8, but the mean correlation for the BCGs is again offset, to higher velocity dispersion at a given CO$_{EW}$. The mean CO$_{EW}$ is almost identical for the two samples, whilst the BCGs have a much higher average velocity dispersion, and hence mass, as expected.  \\begin{figure} \\centerline{\\psfig{figure=bcg_fig7.eps,width=0.5\\textwidth,angle=270}} \\caption{Mg$_2$ index versus CO Equivalent Width for the Brightest Cluster Galaxies (circles) and Coma cluster ellipticals (crosses). } \\end{figure}  \\begin{figure} \\centerline{\\psfig{figure=bcg_fig8.eps,width=0.5\\textwidth,angle=270}} \\caption{Log of velocity dispersion (kms$^{-1}$) versus CO Equivalent Width for the Brightest Cluster Galaxies (circles) and Coma cluster ellipticals (crosses). } \\end{figure} ",
        "conclusions": "We find that BCGs are much more homogeneous in evolved red stellar content than ellipticals overall, and BCGs are somewhat more homogeneous than Coma cluster ellipticals.  The measured scatter in the CO$_{EW}$ indices for BCGs is comparable to the measurement errors.  We interpret this as implying a more uniform and probably earlier star formation history for BCGs than for normal ellipticals.  Metallicity does not appear to be the controlling parameter of CO absorption strength for elliptical galaxies.  Absolute magnitudes, and magnitude residuals relative to the structure parameter relation of Hoessel \\shortcite{ho:80} correlate well with CO absorption depth.  This may imply the presence of an additional intermediate-age population, or a higher metallicity population, in the galaxies which are over-luminous relative to the mean relation. However, this effect is very small and such populations, if present, must be weaker than in most field and cluster ellipticals, given the high degree of homogeneity in BCG CO$_{EW}$ values.  It may be possible to use these correlations to define higher precision distance indicators for BCGs, which could be used for galaxies with redshifts up to 18,000~kms$^{-1}$, as is indicated by the reduced scatter in R-band absolute magnitude, and the reduced scatter about the structure parameter relation after correction for the correlation with CO$_{EW}$, discussed in section 3.  Recent ROSAT observations reveal that a significant number of the Lauer \\& Postman \\shortcite{la:94} galaxies do not lie at the X-ray centroids of their clusters (Paul Lynam, private communication), and there may be better candidates for the dominant central cluster galaxy. Thus our conclusions may refer more generally to bright galaxies towards cluster centres than to individual BCGs inhabiting the very centre of the cluster potential."
    },
    "0002/astro-ph0002170_arXiv.txt": {
        "abstract": "We present a simple method for evaluating the nonlinear biasing function of galaxies from a redshift survey.  The nonlinear biasing is characterized by the conditional mean of the galaxy density fluctuation given the underlying mass density fluctuation $\\coav$, or by the associated parameters of mean biasing $\\bh$ and nonlinearity $\\bt$ (following Dekel \\& Lahav 1999).  Using the distribution of galaxies in cosmological simulations, at smoothing of a few Mpc, we find that $\\coav$ can be recovered to a good accuracy from the cumulative distribution functions of galaxies and mass, $\\cg(\\delg)$ and $\\cm(\\delm)$, despite the biasing scatter. Then, using a suite of simulations of different cosmological models, we demonstrate that $\\cm(\\delm)$ can be approximated in the mildly nonlinear regime by a cumulative log-normal distribution of $1+\\delm$ with a single parameter $\\sigm$, with deviations that are small compared to the difference between $\\cg$ and $\\cm$. Finally, we show how the nonlinear biasing function can be obtained with adequate accuracy directly from the observed $\\cg$ in redshift space. Thus, the biasing function can be obtained from counts in cells once the rms mass fluctuation at the appropriate scale is assumed a priori. The relative biasing function between different galaxy types is measurable in a similar way. The main source of error is sparse sampling, which requires that the mean galaxy separation be smaller than the smoothing scale. Once applied to redshift surveys such as PSC$z$, 2dF, SDSS, or DEEP, the biasing function can provide valuable constraints on galaxy formation and structure evolution. ",
        "introduction": "\\label{sec:intro}  The fact that galaxies of different types cluster differently (\\eg, Dressler 1980; Lahav, Nemiroff \\& Piran 1990; Santiago \\& Strauss 1992; Loveday \\etal 1995; Hermit \\etal 1996; Guzzo \\etal 1997) indicates that the galaxy distribution is in general biased compared to the underlying mass distribution. Cosmological simulations confirm that halos and galaxies must be biased (\\eg, Cen \\& Ostriker 1992; Kauffmann, Nusser \\& Steinmetz 1997; Blanton \\etal 1999; Somerville \\etal 2000). The biasing becomes even more pronounced at high redshift, as predicted by theory (\\eg, Kaiser 1986; Davis \\etal 1985; Bardeen \\etal 1986; Dekel \\& Rees 1987; Mo \\& White 1996; Bagla 1998; Jing \\& Suto 1998; Wechsler \\etal 1998), and confirmed by the strong clustering of galaxies observed at $z\\sim 3$ (Steidel \\etal 1996; 1998). Knowing the biasing scheme is crucial for extracting dynamical information and cosmological constants from the observed galaxy distribution, and may also be very useful for understanding the process and history of galaxy formation.  The simplest possible biasing model relating the density fluctuation fields of matter and galaxies, $\\delm$ and $\\delg$, is the deterministic and linear relation, $\\delg(\\vx)=b\\,\\delta(\\vx)$, where $b$ is a constant linear biasing parameter. However, this is at best a crude approximation, because it is not self-consistent (\\eg, it does not prevent $\\delg$ from becoming smaller than $-1$ when $b>1$) and is not preserved in time. At any given time, scale and galaxy type, the biasing is expected in general to be nonlinear, i.e., $b$ should vary as a function of $\\delta$. The nonlinearity of dark-matter halo biasing (as well as its dependence on scale, mass and time) is approximated fairly well by the model of Mo \\& White (1996), based on the extended Press-Schechter formalism (Bond \\etal 1991). Improved approximations have been proposed by Jing (1998), Catelan \\etal (1998), Sheth \\& Tormen (1999) and Porciani \\etal (1999). It is quantified further for halos and galaxies using cosmological $N$-body simulations with semi-analytic galaxy formation (\\eg, Somerville \\etal 2000). The biasing is also expected, in general, to be stochastic, in the sense that a range of values of $\\delg$ is possible for any given value of $\\delm$. For example, if the biasing is nonlinear on one scale, it should be different and non-deterministic on any other scale. The origin of the scatter is shot noise as well as the influence of physical quantities other than mass density (\\eg, velocity dispersion, the dimensionality of the local deformation tensor which affects the shape of the collapsing object, etc.) on the efficiency of galaxy formation.  Dekel \\& Lahav (1999) have proposed a general formalism for galaxy biasing, that separates nonlinearity and stochasticity in a natural way. The density fields are treated as random fields, and the biasing is fully characterized by the conditional probability distribution function $P(\\delg\\vert\\delm)$. The constant linear biasing factor $b$ is replaced by a mean {\\it biasing function}, \\be \\coav\\equiv b(\\delm)\\,\\delm , \\label{eq:cond_def} \\ee which can in principle take a wide range of functional forms, restricted by definition to have $\\av{\\delg}=0$ and $\\coav\\geq -1$ for any $\\delm$. The stochasticity is expressed by the higher moments about this mean, such as the conditional variance \\be \\sigb ^2(\\del) \\equiv \\av{\\epsilon^2 |\\delm} /\\sigma^2 , \\quad \\epsilon \\equiv \\delg-\\av{\\delg|\\delm} \\ , \\ee scaled for convenience by the variance of mass fluctuations, $\\sigma^2\\equiv\\av{\\delm^2}$. To second order, the biasing function $b(\\del)$ can be characterized by two parameters: the moments $\\bh$ and $\\bt$, \\be \\bh \\equiv\\ \\av{b(\\del)\\, \\del^2} /\\sig^2 \\quad {\\rm and} \\quad \\bt^2 \\equiv\\ \\av{b^2(\\del)\\, \\del^2} /\\sig^2 \\ . \\ee The parameter $\\bh$ is the natural extension of the linear biasing parameter, measuring the slope of the linear regression of $\\delg$ on $\\delm$, and $\\bt/\\bh$ is a useful measure of non-linearity. The stochasticity is characterized independently by a third parameter, $\\sigb ^2 \\equiv \\av{\\epsilon^2}/\\sig^2$. As has been partly explored by Dekel \\& Lahav (1999), these parameters should enter any nonlinear analysis aimed at extracting the cosmological density parameter $\\Omega$ from a galaxy redshift survey, and are therefore important to measure.  In this paper we propose a simple method to measure the biasing function $b(\\delta)$ and the associated parameters $\\bh$ and $\\bt$ from observed data that are either already available, such as the PSC$z$ redshift survey (Saunders \\etal 2000), or that will soon become available, such as the redshift surveys of 2dF (Colless 1999) and SDSS (\\eg, Loveday \\etal 1998) and high-redshift surveys such as DEEP (Davis \\& Faber 1999). Alternative methods have been proposed to measure the biasing function, using the cumulant correlators of the observed distribution of galaxies in redshift surveys (Szapudi 1998) or their bispectrum (Matarrese, Verde, Heavens 1997, Verde \\etal 1998).  We first show in \\se{CDF}, using halos and galaxies in $N$-body simulations, that the difference between the cumulative distribution functions (CDFs) of galaxies and mass can be straightforwardly translated into $\\coav$ despite the scatter in the biasing scheme. Then, in \\se{rob}, we demonstrate that for our purpose, $\\cm(\\delm)$ is insensitive to the cosmological model and can be approximated robustly by a cumulative log-normal distribution. This means that we do not need to observe $\\cm(\\delta)$, which is hard to do; we only need to measure $\\cg(\\delg)$ and, independently, the rms value $\\sigm$ of the mass fluctuations on the same scale. In \\se{redshift}, we slightly modify the method to account for redshift-space distortions, and use mock galaxy catalogs from N-body simulations to evaluate the associated errors. Finally, in \\se{errors}, we estimate the errors due to the sparse sampling and finite volume. The method and its applications to existing and future data are discussed in \\se{conc}. ",
        "conclusions": "\\label{sec:conc}  We propose a simple prescription for recovering the mean nonlinear biasing function from a large redshift survey. The biasing function is defined by $b(\\delm)\\,\\delm = \\av{\\delg | \\delm}$, and is characterized to second order by two parameters, $\\bh$ and $\\bt$, measuring the mean biasing and its nonlinearity respectively. The method is applied at a given cosmology, time, object type and smoothing scale, and involves one parameter that should be assumed a priori --- the rms mass density fluctuation $\\sigma$ on the relevant scale.  The main steps of the algorithm are as follows: \\begin{enumerate} \\item Obtain the observed cumulative distribution function in redshift space $\\cgz(\\delgz)$, by counts in cells or with window smoothing at a certain smoothing length. \\item Assume a value for $\\sigm$ on that scale and for the cosmological density parameter $\\om$, and approximate the mass CDF in redshift space by $\\clnz(\\delmz;\\sigmz)$, the cumulative log-normal distribution (\\eqd{cln}), with the width $\\sigmz$ derived from $\\sigm$ and $\\om$ by \\equ{sigmz}. \\item Derive the mean biasing function by \\be \\delg (\\delm\\!=\\!\\delmz) \\simeq \\delgz (\\delmz ) = \\cgz^{-1} [\\clnz(\\delmz;\\sigmz)] \\ \\ . \\ee \\end{enumerate}  We first showed that the mean biasing function, at TH8 smoothing, can be derived with reasonable accuracy from the r-space CDFs of galaxies (or halos) and mass, despite the biasing scatter. We then demonstrated that for a wide range of CDM cosmologies the mass CDF can be properly approximated for this purpose by a log-normal distribution of the same width $\\sigm$. Next we showed that the biasing functions in z-space and r-space are very similar, and that the z-space mass CDF can also be approximated by a log-normal distribution, with a width derived from $\\sigm$ via \\equ{sigmz}. This allows us to apply the method directly to the observed CDF in a redshift survey. The errors in the recovered biasing function and its moments, in an ideal case of dense sampling in a large volume, are at the level of a few percent.  In any realistic galaxy survey the limited volume and discrete sampling introduce further random and systematic errors. For a survey like the PSC$z$ survey, the main source of error is the sampling density; the error does not exceed $\\sim 10\\%$ as long as the mean observed galaxy separation is kept smaller than the smoothing radius. We are currently in the process of applying this method to the PSC$z$ survey (E. Branchini, \\etal 2000, in preparation), where a more specific error analysis will be carried out. The sampling errors are expected to be significantly smaller for the upcoming 2dF and SDSS redshift surveys.  In \\se{CDF} we showed that our method works well both for halos and for galaxies, on scales 5 to 15$\\hmpc$, and in the redshift range $0\\leq z \\leq 3$ over which the biasing is expected to change drastically. We obtain a similar accuracy when we vary the cosmological model, the mass of the halos in the comparison, or galaxy properties such as morphological type and luminosity. The approximation $\\delg(\\delta)$ is consistent (the deviation is less than 1-$\\sigma$) with the true average biasing function $\\coav$ over a wide range of $\\delm$ values, which covers 98 -- 99\\% of the volume, depending on redshift and the type of biased objects. This allows us to estimate the moments of the biasing function to within a few percent (see Table~1). The moments of the biasing function are derived from 99.9\\% of the volume (99\\% at $z$=3 and for relative biasing).  The method requires as external parameters the rms mass-density fluctuation $\\sigm$ and the cosmological parameter $\\om$. These can be obtained by joint analyses of constraints from several observational data sets, such as the cluster abundance (\\eg, Eke \\etal 1998), peculiar velocities (\\eg, Dekel \\& Rees 1994; Zaroubi \\etal 1997; Freudling \\etal 1999), CMB anisotropies (\\eg, de Bernardis \\etal 1999), and type Ia supernovae (Riess \\etal 1998; Perlmutter \\etal 1999). Examples for such joint analyses are Bahcall \\etal (1999) and Bridle \\etal (1999).  The method is clearly applicable at $z\\simeq0$ with available redshift surveys and especially with those that will become available in the near future, 2dF and SDSS. In the future, this method may become applicable at higher redshifts as well, where the biasing plays an even more important role. With the accumulation of Lyman-break galaxies at $z\\sim 3$, it may soon become feasible to reconstruct their PDF by counts in cells, and our method will allow a recovery of the biasing function at this early epoch, with consequences on galaxy formation and on the evolution of structure.  We have concentrated here on smoothing scales relevant to galaxy biasing, but the method may also be applicable for the biasing of galaxy clusters, on scales of a few tens of Mpc. The biasing scatter may be larger for clusters because of their sparse sampling, but the larger mean biasing parameter for clusters may help in regaining the required monotonicity for \\equ{cc} to provide a valid approximation to the mean biasing function. The mass PDF has been checked to be properly approximated by a log-normal distribution at smoothing scales in the range 20 to $40\\hmpc$, using simulations of the standard CDM and Cold+Hot DM models (Borgani \\etal 1995). The errors due to sparse sampling would require a smoothing scale at the high end of this range.  In a large redshift survey which distinguishes between object types, one can measure the {\\it relative} biasing function between two object types by applying \\equ{cc12} in redshift space, using the observed CDFs for the two types without appealing to the underlying mass distribution at all.  The upcoming large redshift surveys 2dF and SDSS, and the DEEP survey at $z\\sim 1$, are indeed expected to provide adequate samples of different galaxy types. Compared with the predictions of simulations and semi-analytic modeling of galaxy formation (\\eg, Kauffmann \\etal 1999; Benson \\etal 1998; Baugh \\etal 1999; Somerville \\& Primack 1999), the measured relative biasing function can provide valuable constraints on the formation of galaxies and the evolution of structure.  While implementing the method outlined above for measuring the mean nonlinear biasing function using current and future redshift surveys, the next challenge is to devise a practical method for measuring the biasing scatter about the mean."
    },
    "0002/astro-ph0002346_arXiv.txt": {
        "abstract": "The systematic search for the presence of cyclotron lines in the spectra of accreting X-ray pulsars is being carried on with the BeppoSAX satellite since the beginning of the mission. These highly successful observations allowed the detection of cyclotron lines in many of the accreting X-ray pulsars observed. Some correlations between the different measured parameters were found. We present these correlations and discuss them in the framework of the current theoretical scenario for the X--ray emission from these sources. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002420_arXiv.txt": {
        "abstract": "We present the first results from the Shellflow program, an all-sky Tully-Fisher (TF) peculiar velocity survey of 276 Sb$-$Sc galaxies with redshifts between 4500 and 7000 \\kms. Shellflow was designed to minimize systematic errors between observing runs and between telescopes, thereby removing the possibility of a spurious bulk flow caused by data inhomogeneity. A fit to the data yields a bulk flow amplitude $V_{{\\rm bulk}} = 70^{+100}_{-70}\\ \\kms$ ($1\\sigma$ error) with respect to the Cosmic Microwave Background, i.e., consistent with being at rest. At the 95\\% confidence level, the flow amplitude is $< 300\\,\\kms.$ Our results are insensitive to which Galactic extinction maps we use, and to the parameterization of the TF relation. The larger bulk motion found in analyses of the Mark III peculiar velocity catalog are thus likely to be due to non-uniformities between the subsamples making up Mark III. The absence of bulk flow is consistent with the study of Giovanelli and collaborators and flow field predictions from the observed distribution of IRAS galaxies. ",
        "introduction": "\\setcounter{footnote}{0}  It is of great cosmological importance to identify the volume of space, centered on the Local Group, which is at rest with respect to the Cosmic Microwave Background radiation (CMB). Very large-scale density fluctuations are required to move large volumes of galaxies in the gravitational instability picture of structure formation. In standard Cold Dark Matter (CDM) cosmogonies, density fluctuations on scales $\\simgt 100\\hmpc$ are very small. As a result, the volume of space encompassed by the nearest superclusters (Great Attractor, Pisces-Perseus, Coma) is expected to be nearly at rest with respect to the CMB, and the distribution of matter {\\em within\\/} this volume should explain the $\\sim 600$ \\kms\\ motion of the Local Group in the CMB frame. However, the detection of a large amplitude flow ($V_{{\\rm bulk}} \\simgt 700$ \\kms) out to 15,000 \\kms\\ by Lauer \\& Postman (1994), along with recent measurements of similar amplitude (although different directions) by Willick (1999b) and Hudson \\etal\\ (1999), have challenged the notion that the bulk flow on large scales is small, and are pushing CDM models to the breaking point (e.g., Feldman \\& Watkins 1994; Strauss \\etal\\ 1995).  However, Giovanelli \\etal\\ (1998a,b) and Dale \\etal\\ (1999) find results consistent with no flow in their survey of field and cluster spirals out to 20,000 \\kms.  The measured bulk flow on smaller scales is also controversial. The most recent POTENT reconstructions (Dekel \\etal\\ 1999) of the Mark III velocities (Willick \\etal\\ 1997; Mark III) find a bulk velocity within 6000 \\kms\\ of $370 \\pm 110$ \\kms\\ in the CMB frame towards Supergalactic $(L,B)=(165^\\circ,\\, -10^\\circ)$\\footnote{A slightly smaller bulk flow amplitude of $305 \\pm 110\\kms$ is obtained if the VELMOD2 TF calibration of Willick \\& Strauss (1998) is used.}. Dekel \\etal\\ (1999) argue that this motion is generated by the {\\it external} mass distribution on very large scales (see also Courteau \\etal\\ 1993).  On the other hand, Giovanelli \\etal\\ (1998) find a flow consistent with zero on similar scales from their field sample, a result consistent with the surface brightness fluctuation data of Tonry \\etal\\ (2000) and SN Ia distances (Riess 2000).  Accurate ($\\simlt 150\\ \\kms$) measurement of the bulk flow at 6000 \\kms\\ requires that the galaxy distance data be homogeneous and free of systematic effects at the $2-3$\\% level. This cannot be guaranteed for datasets, such as the Mark III catalog, that are composed of two or more independent peculiar velocity surveys. Indeed, Willick \\& Strauss (1998) found evidence of systematic errors in the relative zero points of the various TF samples that make up the Mark III catalog. Thus, the controversy over the observed bulk flow within $60 \\hmpc$ stems, in large part, from the difficulty of combining the various galaxy distance samples used in flow studies into a single homogeneous catalog. None of the previous surveys extending to $60\\hmpc$ sampled the {\\it entire} sky uniformly and reduced the raw data for Northern and Southern hemisphere galaxies using identical techniques\\footnote{Earlier attempts include Roth (1994) and Schlegel (1995). The work of Giovanelli \\etal\\ (1998a,b) incorporates the Southern galaxy survey of Mathewson \\etal\\ (1992), but these authors claim to have reduced the systematic offset in the calibration between the two data sets to negligible levels.  On scales larger than 6000 \\kms, the surveys of Lauer \\& Postman (1994) and Dale \\etal\\ (1999) were designed in a manner analogous to Shellflow.}.  To address these issues we undertook a new TF survey focussed on a relatively narrow redshift shell centered at $\\sim 6000\\ \\kms.$ Our survey, ``Shellflow,'' was designed to provide {\\em precise\\/} and {\\em uniform\\/} photometric and spectroscopic data over the whole sky, and thus to remove the uncertainties associated with matching heterogeneous data sets. In this {\\em Letter,} we report the first scientific result from Shellflow: a reliable, high-accuracy measurement of the bulk flow at $60\\hmpc.$ In future papers (Willick \\etal\\ 2000, Paper II; Courteau \\etal\\ 2000, Paper III), we will describe the data set in greater detail and address related scientific questions, including higher-order moments of the flow field and the value of $\\beta\\equiv \\Omega_m^{0.6}/b.$ ",
        "conclusions": "The results we have presented here are in broad agreement with other recently reported results on the flow field in the local universe.  These include the analyses of the SCI and SFI TF samples (Giovanelli \\etal\\ 1998a,b), who find V$_{\\rm bulk}=200\\pm65$ \\kms\\ within 6500 \\kms\\ and no motion for shells farther than 5000 \\kms; a similar analysis by Dale et al. (1999) who find no significant motion of clusters between 5000 and 20000 \\kms; as well as work from Tonry \\etal\\ (2000), who obtain V$_{\\rm bulk}=289\\pm137$ \\kms\\ at 3000 \\kms\\ from surface brightness fluctuation data, and Riess (2000), who finds no measurable bulk flow in the CMB frame from a sample of 44 SNe Ia with an average depth of 6000 \\kms. Taken together these results suggest that by a distance of $60\\hmpc$, we are seeing a convergence of the flow field to the CMB frame, as is predicted by the observed distribution of IRAS galaxies (Strauss \\etal\\ 1992; Schmoldt \\etal\\ 1999; Rowan-Robinson \\etal\\ 2000). While the data for more distant samples remain ambiguous, with several claims of large amplitude flows on scales $\\simgt 100\\hmpc,$ the results within $60\\hmpc$ cast serious doubt on these claims. If, as abundant evidence suggests, the universe monotonically approaches homogeneity on ever larger scales, it is difficult to see how $\\simgt 600\\,\\kms$ bulk flows on $\\simgt 100\\hmpc$ scales can be reconciled with negligible bulk flow on a scale half as large. From this perspective it seems likely that the results of Lauer \\& Postman (1994), Willick (1999b), and Hudson \\etal\\ (1999) are due, at least in part, to subtle and small systematic effects.  In summary, we find no significant motion of a shell of galaxies centered at 6000 \\kms, as seen in the CMB frame. Equivalently, from the vantage point of the LG frame, we see a motion equal in amplitude and opposite in direction to the motion of the LG through the CMB. Our results are insensitive to whether we adopt the BH or the SFD reddenings, as well as to the parameterization of the TF relation. Future papers will present the spectroscopic and photometric data, give a detailed account of our TF analysis, including tests for a surface-brightness dependence of the TF relation, consider higher-order moments of the velocity field, and compare with the IRAS-predicted velocity field, following the methods of Davis, Nusser, \\& Willick (1996) and Willick \\& Strauss (1998), to estimate $\\beta=\\Omega_m^{0.6}/b.$  We will also use the Shellflow sample to recalibrate and homogeneously merge the major TF catalogs out to 6000 \\kms, including Mark III and SFI (Haynes \\etal\\ 1999). Such a future superset of existing TF catalogs, based on a reliable, all-sky calibration, will provide a powerful tool for studying the velocity and density fields in the local universe."
    },
    "0002/astro-ph0002085_arXiv.txt": {
        "abstract": "Naturally occurring water vapor maser emission at 1.35 cm wavelength provides an accurate probe for the study of accretion disks around highly compact objects, thought to be black holes, in the centers of active galaxies. Because of the exceptionally fine angular resolution, 200 microarcseconds, obtainable with very long baseline interferometry, accompanied by high spectral resolution, $<0.1$~\\kms, the dynamics and structures of these disks can be probed with exceptional clarity. The data on the galaxy NGC\\_4258 are discussed here in detail. The mass of the black hole binding the accretion disk is $3.9 \\times 10^7$~\\msun. Although the accretion disk has a rotational period of about 800 years, the physical motions of the masers have been directly measured with VLBI over a period of a few years. These measurements also allow the distance from the earth to the black hole to be estimated to an accuracy of 4 percent. The status of the search for other maser/black hole candidates is also discussed. ",
        "introduction": "The observational evidence for the existence of supermassive black holes ($10^6$--$10^9$ times the mass of the sun, \\msun) in the centers of active galaxies has been accumulating at an ever accelerating pace for the last few decades (e.g., Rees 1998; Blandford \\& Gehrels 1999). Seyfert (1943) first drew attention to a group of galaxies with unusual excitation conditions in their nuclei indicative of energetic activity. Among the twelve galaxies in his list was NGC\\_4258, which is the subject of much of this paper. Such galaxies, now known as galaxies with active galactic nuclei (AGN), have grown in membership and importance. Ironically, NGC\\_4258  no longer  belongs to the class of Seyfert galaxies  by modern classification standards (Heckman 1980), but it is still considered to have a mildly active galactic nucleus. Meanwhile, the study of AGN has become a major field in modern astrophysics. In the 1960s, galaxies with AGN were discovered with intense radio emission arising from  jets of relativistic particles often extending far beyond the optical boundaries of the host galaxy. The central engine, the source of energy that powers such jets and other phenomena in the centers of galaxies, has long been ascribed to black holes (e.g., Salpeter 1964; Blandford \\& Rees 1992). There are two sources of energy for these phenomena: the gravitational energy from material falling onto the black hole and the spin energy of the black hole itself (Blandford \\& Znajek 1977).  The direct evidence for black holes in AGN has come principally from observations of the motions of gas and stars in the  extended environments of black holes. In the optical and infrared domains, the evidence for black holes from stellar measurements comes from an analysis of the velocity dispersion of stars as a function of distance from the dynamical centers of galaxies. In the case of our own Galactic center, the proper motions (angular velocities in the plane of the sky) of individual stars can be measured. These data show that there is a mass of about $2.6 \\times 10^6$~\\Mdot\\ within a volume of radius 0.01 pc (Genzel et al. 1997; Ghez et al. 1998). In addition, measurements by the Hubble Space Telescope of the velocity field of hydrogen gas in active galaxies indicate the presence of  massive centrally condensed objects. Reviews of these data have been written by Faber (1999), Ho (1999), Kormendy and Richstone (1995), and others.  In the X-ray portion of the spectrum, there is compelling evidence for black holes in AGN from the detection of the highly broadened iron K$\\alpha$ line at 6.4 keV. The line is broadened by the gravitational redshift of gas as close as 3 Schwarzschild radii from the black hole. An example of an iron line profile in the galaxy MCG\\_-6-30-15 is shown in Figure~1 (Tanaka et al. 1995). The linear extent of the emission region cannot be determined directly by the X-ray telescope, so it is not possible to estimate directly the mass of the putative black hole. Detailed analysis of the line profile suggests that the black hole is spinning (e.g., Bromley, Miller, \\& Pariev 1998).  \\begin{figure}[t] \\plotfiddle{fig1.eps}{3.25in}{.2}{90}{90}{-275}{-204} \\caption{\\setlength{\\baselineskip}{8 pt} The X-ray spectrum of the galaxy MCG\\_-6-30-15, observed by the Japanese ASCA satellite. The top panel shows the total spectrum with a model of the continuum emission fitted to the data outside the range of 5--7 keV. The bottom panel shows the residuals, which reveal a broad spectral feature attributed to the Fe K$\\alpha$ line at 6.4 keV. The line has a width of 100,000 \\kms. The most extreme redshifted part is thought to arise from gas at a radius of about 3 Schwarzschild radii. (From Tanaka et al. 1995)} \\end{figure}   In the radio regime, a new line of inquiry has given unexpectedly clear and compelling evidence for black holes: the discovery of water masers orbiting highly massive and compact central objects. With the aid of very long baseline interferometry (VLBI), which provides angular resolution as fine as 200~microarcseconds ($\\mu$as) at a wavelength of 1.3~cm and spectral resolution of 0.1 km$\\,$s$^{-1}$ or less, the structure of accreting material around these central objects can be studied in detail. This paper describes the observations and the significance of these measurements of water masers in AGN. We begin with a brief description of cosmic masers and the interferometric techniques used to observe them. ",
        "conclusions": ""
    },
    "0002/astro-ph0002058_arXiv.txt": {
        "abstract": "The nature of the dark matter in the haloes of galaxies is one of the outstanding questions in astrophysics.  All stellar candidates, until recently thought to be likely baryonic contributions to the Halo of our Galaxy, are shown to be ruled out.  Faint stars and brown dwarfs are found to constitute only a few percent of the mass of the Galaxy. Stellar remnants, including white dwarfs and neutron stars, are shown to be very constrained as well.  High energy gamma-rays observed in HEGRA data place the strongest constraints, $\\Omega_{WD} < 3 \\times 10^{-3} h^{-1}$, where $h$ is the Hubble constant in units of 100 km s$^{-1}$ Mpc$^{-1}$.  Hence one is left with several unanswered questions: 1) What are MACHOs seen in microlensing surveys? 2) What is the dark matter in our Galaxy?  Indeed a nonbaryonic component in the Halo seems to be required. ",
        "introduction": "The nature of the dark matter in the haloes of galaxies is an outstanding problem in astrophysics.  Over the last several decades there has been great debate about whether this matter is baryonic or must be exotic.  Many astronomers believed that a stellar or substellar solution to this problem might be the most simple and therefore most plausible explanation. However, in the last few years, these candidates have been ruled out as significant components of the Galactic Halo. I will discuss limits on these stellar candidates, and argue for my personal conviction that: {\\em Most of the dark matter in the Galactic Halo must be nonbaryonic.}  Until recently, stellar candidates for the dark matter, including faint stars, brown dwarfs, white dwarfs, and neutron stars, were extremely popular.  However, recent analysis of various data sets has shown that faint stars and brown dwarfs probably constitute no more than a few percent of the mass of our Galaxy \\cite{kfrbfgk,kfrgf96a,kfrgf96b,kfrmcs,kfrfgb,kfrfreese}. Specifically, using Hubble Space Telescope and parallax data, we showed that faint stars and brown dwarfs contribute no more than 1\\% of the mass density of the Galaxy.  Microlensing experiments (the MACHO \\cite{kfrmacho:1yr}, \\cite{kfrmacho:2yr} and EROS \\cite{kfransari}) experiments), which were designed to look for Massive Compact Halo Objects (MACHOs), also failed to find these light stellar objects and ruled out substellar dark matter candidates in the $(10^{-7} - 10^{-2}) M_\\odot$ mass range.  Recently white dwarfs have received attention as possible dark matter candidates.  Interest in white dwarfs has been motivated by microlensing events interpreted as being in the Halo, with a best fit mass of $\\sim 0.5 M_\\odot $. However, I will show that stellar remnants including white dwarfs and neutron stars are extremely problematic as dark matter candidates. A combination of excessive infrared radiation, mass budget issues and chemical abundances constrains the abundance of stellar remnants in the Halo quite severely, as shown below. Hence, white dwarfs, brown dwarfs, faint stars, and neutron stars are either ruled out or extremely problematic as dark matter candidates.  Thus the puzzle remains, What are the 14 MACHO events that have been interpreted as being in the Halo of the Galaxy?  Are some of them actually located elsewhere, such as in the LMC itself? These questions are currently unanswered.  As regards the dark matter in the Halo of our Galaxy, one is driven to nonbaryonic constituents as the bulk of the matter. Possibilities include supersymmetric particles, axions, primordial black holes, or other exotic candidates.  In this talk I will focus on the arguments against stellar remnants as candidates for a substantial fraction of the dark matter, as white dwarfs in particular have been the focus of attention as potential explanations of microlensing data.  For a discussion of limits on faint stars and brown dwarfs, see earlier conference proceedings by Freese, Fields, and Graff (\\cite{kfrfreese} and \\cite{kfrconf1}). ",
        "conclusions": "\\paragraph{A Zero Macho Halo?} The possibility exists that the 14 microlensing events that have been interpreted as being in the Halo of the Galaxy are in fact due to some other lensing population.  One of the most difficult aspects of microlensing is the degeneracy of the interpretation of the data, so that it is currently impossible to determine whether the lenses lie in the Galactic Halo, or in the Disk of the Milky Way, or in the LMC.  In particular, it is possible that the LMC is thicker than previously thought so that the observed events are due to self-lensing of the LMC.  All these possibilities are being investigated.  More data are required in order to identify where the lenses are.  Microlensing experiments have ruled out baryonic dark matter objects in the mass range $10^{-7}M_\\odot $ all the way up to $10^{-2}M_\\odot$. In this talk I discussed the heavier possibilities in the range $10^{-2}M_\\odot $ to a few $M_\\odot $.  Brown dwarfs and faint stars are ruled out as significant dark matter components; they contribute no more than 1\\% of the Halo mass density.  Stellar remnants are not able to explain the dark matter of the Galaxy either; none of the expected signatures of stellar remnants, i.e., infrared radiation, large baryonic mass budget, and C,N, and He$^4$ abundances, are found observationally.  Hence, in conclusion, \\hfill\\break 1) Nonbaryonic dark matter in our Galaxy seems to be required, and \\hfill\\break 2) The nature of the Machos seen in microlensing experiments and interpreted as the dark matter in the Halo of our Galaxy remains a mystery.  Are we driven to primordial black holes \\cite{kfrcarr} \\cite{kfrjedam}, nonbaryonic Machos (Machismos?), mirror matter Machos (\\cite{kfrmohap}) or perhaps a no-Macho Halo?"
    },
    "0002/astro-ph0002328_arXiv.txt": {
        "abstract": "In this paper we describe the Bayesian link between the cosmological mass function and the distribution of times at which isolated halos of a given mass exist. By assuming that clumps of dark matter undergo monotonic growth on the time-scales of interest, this distribution of times is also the distribution of `creation' times of the halos. This monotonic growth is an inevitable aspect of gravitational instability. The spherical top-hat collapse model is used to estimate the rate at which clumps of dark matter collapse. This gives the prior for the creation time given no information about halo mass. Applying Bayes' theorem then allows {\\em any} mass function to be converted into a distribution of times at which halos of a given mass are created. This general result covers both Gaussian and non-Gaussian models. We also demonstrate how the mass function and the creation time distribution can be combined to give a joint density function, and discuss the relation between the time distribution of major merger events and the formula calculated. Finally, we determine the creation time of halos within three N-body simulations, and compare the link between the mass function and creation rate with the analytic theory. ",
        "introduction": "\\label{sec:intro}  The hierarchical build-up of self-gravitating dark matter is thought to drive evolution in the observable universe. The formation of clumps of dark matter precipitates the formation of galaxies by providing a potential well into which gas can fall and subsequently cool. Violent mergers between equally sized halos and their associated galaxies are thought to be important for starbursts and quasar activation. In order to model and understand the observable universe it is therefore essential to understand the build-up of the dark structure.  The most widely used analytic model for the distribution of mass in isolated halos at any epoch comes from Press-Schechter (PS) theory \\cite{ps}.  By smoothing the initial field of density fluctuations on different scales, information on the distribution of perturbation sizes can be obtained. Linking the time at which these perturbations collapse to the initial overdensities using the simplified spherical top-hat collapse model allows the distribution of mass in isolated halos at any epoch to be determined \\cite{ps,peacock,bond}.  In Percival \\& Miller \\shortcite{ev1} (hereafter paper~I), we used the tenets of PS theory to model the related, but distinct problem of determining the distribution of times at which halos of a given mass are created. Here, `creation' is defined as the epoch at which non-linear collapse is predicted. Two derivations were given, one of which directly used the trajectories invoked in PS theory \\cite{peacock,bond}, and one of which used Bayes' theorem to convert from the PS mass function to a time distribution. The second derivation required the prior for the creation time which was calculated by examining the trajectories model.  In this paper we extend the Bayesian link between the mass function and the creation time distribution to cover any mass function. This is important, not only because it is known that standard PS theory is wrong in detail (e.g. Sheth \\& Tormen 1999), but especially because the new extension applies to mass functions derived from more general density fields including non-Gaussian models (e.g. Matarrese, Verde \\& Jimenez 2000).  First, we adopt the assumption that all clumps monotonically increase in mass on the cosmological time scales of interest. This monotonic growth is an inevitable aspect of gravitational instability. Every epoch should now be thought of as a creation time for a given clump, and we need not make the distinction between the creation time distribution and the distribution of times at which a given halo exists.  In order to convert from a mass function to a distribution in time we require the prior for the creation time. This is the rate at which creation events occur, given no information about the halo mass. In this work we use the spherical top-hat collapse (STHC) model to provide a simple mechanism for determining the required rate. In Section~\\ref{sec:tophat} we derive the link between collapse time and the overdensity at an early epoch for the STHC model within any Friedmann cosmology. Having determined that this relation is independent of halo mass, this leads directly to an approximation to the prior for the creation time, described in Section~\\ref{sec:time}. This is the second major assumption adopted in this paper: that the prior for the creation time is well approximated by this simple model for the break-away of structure from linear expansion. This means that following the two simple assumptions detailed above, we are able to convert any mass function to give the distribution of epochs at which halos of a given mass are created.  Simple models of cosmologically evolving phenomena often adopt an important mass range rather than a specific halo mass (e.g. paper~I, Granato \\etal\\ 1999). In order to use the work presented here in these models, the joint distribution of halos in mass and creation time is required. Although calculating the required joint probability is formally impossible because the equations cannot be properly normalised, a formula with the correct shape can be determined and is presented in Section~\\ref{sec:joint}.  So far we have not made a distinction between the slow accretion of mass onto a halo and major mergers between halos. Such a distinction is important because only major mergers are thought to play a vital role in starbursts and quasar activation (see paper~I). The time distribution calculated in this paper determines when halos existed (or were created by any mechanism assuming monotonic clump growth) which is not necessarily equal to the distribution of merger events. This is discussed in Section~\\ref{sec:mergers}.  Finally, we compare the analytic link between the mass function and the creation rate to the results from three numerical simulations of structure formation in different cosmological models. An analytic fit to the mass function as described by Sheth \\& Tormen \\shortcite{sheth} is adopted and is converted into a creation rate using the STHC model. This model is compared with and shown to be in good agreement with the numerical results. ",
        "conclusions": "We have demonstrated a simple method for linking any mass function to the corresponding distribution of times at which isolated halos of a given mass are created. In order to provide this link we adopted the assumption that the time scales of interest are those over which the mass of every clump can be thought of as monotonically increasing. The prior for the collapse time was estimated using the STHC model which ties in directly with PS theory, although the method does not use any of PS theory beyond that of the STHC model. We have presented a new derivation of the link between the collapse time and initial overdensity for this model which explicitly shows that this link is independent of the halo mass and is applicable in any Friedmann cosmology. Multiplying the mass function by a function with no mass dependence and proportional to the time derivative of the critical overdensity then provides a joint density function with the correct behaviour for the creation of a halo in mass {\\em and} time. Integrating over the resulting joint density function will give the correct relative number densities of halos within different mass and time intervals.  We have extended the analysis of N-body simulation results presented in paper~I to cover three simulations of the build-up of dark matter within different cosmological models. Rather than using PS theory, we have demonstrated how a fit to the mass function may be converted to give a creation rate. Out of the three functions we have compared to the mass function data, the best fit model for these data when converted to a creation rate also fits the creation rate data the best. This gives us confidence that the formalism presented here is sound, and should give accurate results in more general situations, in particular non-Gaussian models."
    },
    "0002/astro-ph0002352_arXiv.txt": {
        "abstract": "It has been known for a long time (Mestel~1953) that the meridional circulation velocity in stars, in the presence of $\\displaystyle \\mu $-gradients, is the sum of two terms, one due to the classical thermal imbalance ($\\displaystyle \\Omega$-currents) and the other one due to the induced horizontal $\\displaystyle \\mu $-gradients ($\\displaystyle \\mu $-induced currents, or $\\displaystyle \\mu $-currents in short). In the most general cases, $\\displaystyle \\mu $-currents are opposite to $\\displaystyle \\Omega$-currents.  Vauclair (1999) has shown that such processes can, in specific cases, lead to a quasi-equilibrium stage in which both the circulation and the helium settling is frozen. Here we present computations of the circulation currents in halo star models, along the whole evolutionary sequences for four stellar masses with a metallicity of [Fe/H] = -2. We show that such a self-regulated process can account for the constancy of the lithium abundances and the small dispersion in the Spite plateau. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002164_arXiv.txt": {
        "abstract": "The dominant linear contribution to cosmic microwave background (CMB) fluctuations at small angular scales ($\\lesssim 1'$) is a second-order contribution known as the Vishniac or Ostriker-Vishniac effect.  This effect is caused by the scattering of CMB photons off free electrons after the universe has been reionized, and is dominated by linear perturbations near the $R_V =2$ Mpc/($h \\Gamma/0.2)$ scale in the Cold Dark Matter cosmogony.  As the reionization of the universe requires that nonlinear objects exist on some scale, however, one can compare the scale responsible for reionization to $R_V$ and ask if a linear treatment is even feasible in different scenarios of reionization.  For an $\\Omega_0 = 1$ cosmology normalized to cluster abundances, only $\\sim 65 \\%$ of the linear integral is valid if reionization is due to quasars in halos of mass $\\sim 10^9 M_\\odot$, while $\\sim 75\\%$ of the integral is valid if reionization was caused by stars in halos of $\\sim 10^6 M_\\odot$.  In $\\Lambda$ or open cosmologies, both the redshift of reionization and $z_V$ are pushed further back, but still only $\\sim 75 \\%$ to $\\sim 85 \\%$ of the linear integral is valid, independent of the ionization scenario.   We point out that all odd higher-order moments from Vishniac fluctuations are zero while even moments are non-zero, regardless of the gaussianity of the density perturbations.  This provides a defining characteristic of the Vishniac effect that differentiates it from other secondary perturbations and may be helpful in separating them. ",
        "introduction": "While recombination at $z \\approx 1100$ marked the end of ionized hydrogen from the viewpoint of a linearly evolving universe, the nonlinear evolution of small-scale perturbations resulted in the reionization of the intergalactic medium at much lower redshifts.  The fact that quasar spectra show an absence of an absorption trough from Ly$\\alpha$ resonant scattering by neutral H atoms distributed diffusely along the line of sight, the Gunn-Peterson effect (Gunn \\& Peterson 1965), means that this reionization must have occurred with a high degree of efficiency before a redshift of 5.  One of the necessary consequences of this reionization is the presence of secondary anisotropies in the cosmic microwave background (CMB) due to the scattering of photons off ionized electrons.  These secondary fluctuations can be divided into two classes: anisotropies due to nonlinear structures and linear anisotropies.  Nonlinear secondary anisotropies are of several types.  Some of the more studied of these include the scattering of photons off the hot intracluster medium of galaxy clusters (Sunyaev \\& Zel'dovich 1970, 1972; or for more recent treatments see, e.g., Evrard \\& Henry 1991; Colfrancesco et al.\\ 1994; Aghanim et al.\\ 1997), gravitational lensing (see, e.g., Linder 1997; Metcalf \\& Silk 1997), the impact of inhomogeneous reionization (Aghanim et al.\\ 1995; Peebles \\& Juszkiewicz 1998; Knox, Scoccimarro, \\& Dodelson 1998), and the Rees-Sciama effect due to the bulk motions of collapsing nonlinear structures (see, e.g., Rees \\& Sciama 1968; Kaiser 1982; Seljak 1996).  Small-scale linear anisotropies come in fewer flavors.  Detailed analyses of linear perturbations have uncovered a single dominant effect known as the Vishniac or Ostriker-Vishniac effect (Hu, Scott, \\& Silk 1994; Dodelson \\& Jubas 1995; Hu \\& White 1995; Hu \\& Sugiyama 1996).  The level of these perturbations has been calculated by several authors (Ostriker \\& Vishniac 1985; Vishniac 1987; Jaffe \\& Kamionkowski 1998, hereafter JK).  These investigations raise the question of whether a detectable Vishniac effect even exists since nonlinear structures must exist on {\\em some} length scale at the time of secondary scattering of CMB photons, as it is only by the formation of nonlinear objects that the universe is able to reionize itself.  If these scales are comparable to those making the dominant contribution to the Vishniac effect, then a linear analysis is inappropriate and a calculation of secondary anisotropies must incorporate nonlinear effects.  In this work we determine the minimum length scale, $R_V$, which must remain linear in order for a linear approach to scattering by ionized regions with varying bulk motions to be accurate for the range of angular scales over which one can hope to measure secondary fluctuations. In hierarchical scenarios of structure formation, such as the Cold Dark Matter (CDM) model, smaller structures assemble at early times, later merging to form larger objects.  This allows us to place limits on the time between the formation of structures large enough to reionize the universe and the time at which $R_V$ becomes nonlinear. At that point, while peculiar velocities of ionized gas continue to be imprinted on the microwave background, the nature of this signature is qualitatively different and is best interpreted from another perspective.  The structure of this work is as follows. In Sec.\\ 2 we describe the Vishniac effect and determine the physical length scale on which it depends.  In Sec.\\ 3 we compare this to the scale of reionizing objects in different reionization scenarios and discuss the applicability of linear theory.  In Sec.\\ 4 we examine how the Vishniac effect is distinguished from other effects.  Conclusions are given in Sec.\\ 5, and the various cosmological expressions used throughout are summarized in the appendix. ",
        "conclusions": "Due to the tremendous predictive power of linear theory, comparisons between linear predictions and large-scale cosmic microwave background measurements promise to constrain cosmological parameters to the order of a few percent (Jungman et al.\\ 1996; Bond, Efstathiou, \\& Tegmark 1997; Zaldarriaga, Spergel, \\& Seljak 1997).  The natural extension of this approach is to try to measure small-scale secondary anisotropies and match them to linear predictions to study the reionization history of universe.  The situation in this case is more muddled, however, as a number of nonlinear secondary effects also contribute at these scales.  The dominant secondary linear anisotropy is a second-order contribution known as the Vishniac or Ostriker-Vishniac effect.  As this effect can be predicted accurately as a function of cosmological parameters, several authors have proposed that its measurement will prove to be a sensitive probe of the reionization history of the universe.  Reionization occurs by the formation of nonlinear structures, however, raising the question of whether a regime of redshift space exists in which these objects have collapsed but a linear analysis is still appropriate.  In this work, we have determined the relevant physical scales that give rise to the Vishniac effect in a Cold Dark Matter cosmogony, showing that approximations are already compromised when $1$ Mpc/($h \\Gamma_{0.20}$) scales have become nonlinear, and break down when $2$ Mpc/($h \\Gamma_{0.20}$) dark matter halos reach overdensities of 1. The width of the redshift regime over which the effect can be imprinted on the CMB is dependent on the cosmological parameters and the reionizing mass scale.  Schemes in which reionization is due to radiation from active galactic nuclei associated with dark matter halos of masses $\\gtrsim 10^9 M_\\odot$ are limited by the absence of a Gunn-Peterson absorption trough.  As reionization must have occurred with a high degree of efficiency before a redshift of 5, such models are successful only if one assumes a large value of $\\sigma_8$, or considers open models with slowly-changing linear growth factors. Both these assumptions push back the redshift at which $R_V$ becomes nonlinear, limiting the range over which a linear analysis is appropriate.  Scenarios in which reionization is due to much smaller objects, such as stars formed in dwarf galaxies associated with dark matter halos of masses $\\gtrsim 10^6 M_\\odot$, are able to reionize the universe at much larger redshifts even in cosmologies in which $\\sigma_8$ is small and $D(z)$ quickly evolving.  This represents only a marginal gain however, as the high redshift contribution to the Vishniac integral is roughly proportional to comoving distance, and comoving distances are small at high redshifts.  Thus low-mass scenarios of reionization are more compatible with a linear analysis not so much because they reionize earlier as because they allow $R_V$ to become nonlinear more recently without violating Gunn-Peterson limits.  The Vishniac effect arises from physical processes that are distinct from nonlinear secondary anisotropies.  Its detection indicates the presence of a redshift regime in which a delicate cancellation takes place due to the lack of curl in the peculiar velocity field and slow variations in $G(w)$.  This leaves a unique signature in the higher-order moments of the temperature fluctuations that is absent from its nonlinear counterparts.  Furthermore, due to the predictive power of linear theory, it represents a sensitive probe of the reionization history not available from measurements of nonlinear contributions.  As with measurements of large angular scale anisotropies, small-scale microwave background anisotropy measurements have the potential to uncover much about the history of our universe.  Also as with large-scale measurements, whether this potential will be realized remains to be seen.  While the Vishniac effect represents a possible probe of the reionization epoch, the analysis will, as always, be more involved than first suggested.  Ultimately it will only be through the measurement and analysis of small-scale microwave background anisotropies that we will be able to know if there is a detectable Vishniac effect."
    },
    "0002/astro-ph0002487_arXiv.txt": {
        "abstract": "We report the results of analysis of observations of the Vela Pulsar by PCA on RXTE. Our data consists of two parts. The first part contains observations at 1, 4, and 9 days  after the glitch in 1996 and has 27000 sec.  total exposure time. The second part of observations were performed three months after this glitch and have a total exposure time of 93000 sec. We found pulsations in both sets. The observed spectrum is a power-law with no apparent change in flux or count rate. The theoretical expectations of increase in flux due to internal heating after a glitch are smaller than the uncertainty of the observations. ",
        "introduction": "\\label{velaintro} We present observations of the \\object{Vela pulsar} with the Proportional Counter Array (PCA), on the Rossi X-Ray Timing Explorer (RXTE). Our observations cover two distinct time-spans. The first part is very close to the glitch on 1996 October 13.394 UT (Flanagan \\cite{glidate}). It consists of three observations at one, four and nine days after the glitch. We analyzed these sets of data separately. The second series of observations were obtained in January 1997. All data sets of January 1997 were analyzed together. The exact dates of observations are given in Table~\\ref{the_table}.  We first performed spectral analysis of our data, calculated time averaged flux for different observations and put upper limits for the flux change. Then, by using radio ephemerides, we detected the pulsations in the data, and investigated the changes in pulse shape and pulse fraction. Finally, we compared our results with the theoretical expectations of change in flux which might arise because of glitch induced energy dissipation in the neutron star.  Time averaged spectrum analysis is explained in section \\ref{pavspec}. The detected pulse shapes are presented in section \\ref{velatiming}. In section \\ref{velaconc} we discuss the implications of our results. ",
        "conclusions": "\\label{velaconc} \\subsection{Time Averaged Spectrum} The power-law spectrum observed is in agreement with expectations deduced from previous observations of Vela at higher and lower energies(\\\"Ogelman et al. \\cite{vel_og}, Kanbach et al. \\cite{kanbach}, Strickman et al. \\cite{strick96}, Kuiper et al. \\cite{kuiper} ). At this part of spectrum (2-20 keV), the contribution of the pulsar is very small compared to the contribution of the compact nebula surrounding it. As a result the pulse shapes have a very high DC level, as can be seen in Fig.~\\ref{pulses}.  The slightly higher residuals near 6 keV and lower residuals near 4 keV are not characteristics of observed sources, but are artifacts of PCA. This effect, which is a result of the L edge of Xenon, is reduced by the version of response matrices in use, but not completely removed.  Our main conclusion from the analysis is that the spectrum does not change from early post-glitch to late observations. It is a power-law with an index around 2 for all of the observations. The power-law index does not change significantly among the observations. The highest value calculated for the index is 2.107 and the lowest value is 2.009.  This corresponds to a change of 5\\%, which is a fractional upper limit for the change of power-index during the observations.  The upper and lower limits of the flux calculated by the comparison explained in section \\ref{pavspec} and presented in Table~\\ref{the_table} are well within the range of systematic errors. We therefore adopt the systematic errors as the upper limits to any variation in flux.  There is seemingly a jump in the flux between 2-20 keV, from the first to the second observation. This observation is only four days away from the glitch. The pre-glitch temperature of the surface of the Vela pulsar is thought to be around 0.15 keV(\\\"Ogelman et al. \\cite{vel_og}). Theoretical models (Van Riper et al.\\cite{nsth1}; Umeda et al.\\cite{nsth2}; Hirano et al. \\cite{nsth3}) predict an increase at most by a factor of 8, which brings the temperature to 1.2 keV. Attempts to find a blackbody component in this observation did not give significantly different results from other observations. This suggests that the observed flux changes may have little or nothing to do with changes in surface temperature.  Another possible interpretation is that there is an error in the analysis of this particular observation, possibly arising from the calculation of synthetic background. Vela is a faint source for PCA. An improved model in the estimation of background for faint sources has been released by the PCA Team in 1998. This model has been used throughout the calculations. There may be further improvements on the background models that could change the calculated flux. The presented flux is calculated by using the spectrum model, rather than by direct observation. Apart from this observation there is no apparent change in flux or count rate.  Treating the calculated fluxes as very high upper bounds to the Wien tail of possible blackbody radiation from the neutron star surface could in principle be used to rule out some of the models for the post-glitch thermal emission from the neutron stars. In practice this does not work since the surface temperature range of the Vela pulsar is far below the RXTE-PCA energies. \\subsection{Timing Analysis and Pulse Shapes} The epoch of the second ephemeris taken from the Princeton database, 50379.0 MJD, is pretty close to the third observation (see Fig.~\\ref{obseph}), but using the ephemeris alone for the observation does not give a pulse shape. This may be due to the existence of two distinct decay time scales of Vela, 3 days and 30 days, which were observed in all previous glitches and fall within the ranges of ephemeris (Alpar et al. \\cite{Alpar93}). Our data is not good enough to determine any exponential decay time scales.  In view of the rapidly varying period at those epochs, the pulse shapes of the first part of the observations were obtained by a careful interpolation amounting to the construction of an ephemeris that can represent the rapid changes in the pulsar's timing parameters in this postglitch epoch. The phase difference of the two observed peaks in these shapes is the same as the phase difference in the second set of observations (in January 1997). This gives us some confidence in the resultant pulse shapes.  The pulse shapes obtained are not reliable for drawing conclusions on the changes of pulse shape or pulsed fraction, since both of these factors are sensitively dependent on ephemeris. This is best seen by comparing Fig.~\\ref{pulses}~(b) and (c). They belong to the same set of data but have obvious differences both in the pulse shape and pulsed fraction.  \\subsection{Possible Future Work} Extracting the contribution of the compact nebula from the spectrum may help to delineate effects of temperature changes on the neutron star surface. Although the pulsations are detected, they are not reliable enough to justify taking the off-peak photon counts as background to the peak photon counts to remove the effects coming from the DC signal. The field of view of the PCA detector is one degree (RXTE GOF \\cite{PCA}), consequently the compact nebula surrounding the Vela pulsar has a significant contribution to the observed spectrum. The images showing the emission from the pulsar and the sources around it, in particular the compact nebula, can be found in Markwardt (\\cite{craig_neb}), Frail et al. (\\cite{ogel2}), Harnden et al. (\\cite{1985}), and Willmore et al. (\\cite{willmore}).  When we divided the data from the second part of observations into smaller time intervals we have observed that the pulse shape begins to disappear for data strings covering less than 30000 seconds. The exposure time of the data sets we used in this work are below 10000 seconds. This explains the uncertainty in pulse shapes and fractions. Future target of opportunity observation of the Vela pulsar by RXTE need to be allocated more observation time, and should contain observations made approximately 20 days after the glitch, since this is about the time that the surface temperature will reach its maximum according to theoretical models. Also, more detailed ephemerides fitting the post-glitch behavior of the pulsar is necessary to make deductions on changes in pulse shape.  Finally we note that the question of glitch associated energy dissipation in the Vela pulsar has been addressed also with ROSAT observations. Comparison of observations at epochs before and after the glitch has not yielded stringent constraints on the glitch related energy dissipation (Seward et al. \\cite{ROSAT}).  While this work was in preparation another analysis of RXTE/PCA observations of the Vela pulsar was published by Strickman et al. (\\cite{strick2}). They have also detected a pulsed emission and a power-law spectrum. Our analysis differs from theirs in two ways. Their phase-resolved spectra are obtained by taking ``off-pulse'' photons as background to ``on-pulse'' photons, whereas we calculated only time averaged spectra. Another difference is that these authors used data coming from only the first xenon layer for energies below 8 keV, but included data coming from the other two layers for higher energies. In our analysis we used photons detected only in the first xenon layer. As a result of these differences, their power-law index is smaller than the value that we found."
    },
    "0002/astro-ph0002508_arXiv.txt": {
        "abstract": "Moderate resolution, near-IR spectroscopy of MS1512-cB58 is presented, obtained during commissioning of the the Near IR Spectrometer (NIRSPEC) on the Keck II telescope.  The strong lensing of this $z=2.72$~galaxy by the foreground cluster MS1512+36 makes it the best candidate for detailed study of the rest-frame optical properties of Lyman Break Galaxies.  In eighty minutes of on-source integration, we have detected \\ha, \\nii $\\lambda$6583,6548\\AA, \\oi$\\lambda$6300\\AA, He I $\\lambda$5876\\AA, \\oiii $\\lambda$5007,4959\\AA, \\hb, \\hgamma, \\oii$\\lambda$3727, and a strong continuum signal in the range 1.29-2.46\\mic.  A redshift of $z=2.7290\\pm 0.0007$~is inferred from the emission lines, in contrast to the $z=2.7233$~calculated from UV observations of interstellar absorption lines.  Using the Balmer line ratios, we find an extinction of \\ebmv=0.27.  Using the line strengths, we infer an SFR$=620\\pm 18$~\\Myr (\\h0=75, \\q0=0.1, $\\Lambda =0$), a factor of 2 higher than that measured from narrow-band imaging observations of the galaxy, but a factor of almost 4 lower than the SFR inferred from the UV continuum luminosity.  The width of the Balmer lines yields a mass of $M_{vir}=1.2\\times 10^{10}$~\\Msun.  We find that the oxygen abundance is 1/3 solar, in good agreement with other estimates of the metallicity.  However, we infer a high nitrogen abundance, which may argue for the presence of an older stellar population. ",
        "introduction": "The largest sample of high redshift galaxies was selected from observations of the extinction of rest-frame far-ultraviolet light by intrinsic and intergalactic absorption.  These Lyman Break galaxies (hereafter LBGs; Steidel et al. 1996) may, in principle, be the tracers of the global star formation history of the universe (Madau et al. 1998).  However, the unquantified effects of dust extinction present an obstacle to such interpretation.  For example, corrections to the star formation rate (SFR) of LBGs have been suggested to be as large as factors of 2--10 (Pettini et al.  1998, hereafter P98; Trager et al.  1997).  In this paper we present the first observations of an LBG with NIRSPEC, the near infrared spectrometer on the Keck II telescope. We will use the rest-frame optical spectrum to obtain a model independent measure of dust extinction, an estimate of the virial mass, and the star formation rate. Spectra such as this one also provide the best method for determining the metal abundance in LBGs.  We will examine the evolutionary status of cB58, and its implications as a ``typical'' LBG.  We have chosen the brightest LBG as our first target (MS1512-cB58; see Yee et al. 1996).  MS1512-cB58 is a gravitationally lensed starburst galaxy at z=2.72.  It was discovered serendipitously, within the field of the z=0.37 galaxy cluster MS1512+36, in spectra obtained for the CNOC survey (Yee, Ellingson, \\& Carlberg 1996).  MS1512-cB58 is clearly extended with the morphology of a lensed arc (Williams \\& Lewis, 1997; Seitz et al., 1998).  It has an ultraviolet spectrum representative of LBGs but due to a factor of $\\sim 30$~magnification (Seitz et al.  1998) it is orders of magnitude brighter than any other object of its kind (V=20.6, \\kp=17.8; Ellingson et al.  1996, hereafter E96).  Like most LBGs, its UV spectrum (devoid of strong, high ionization emission lines) rules out any non-stellar (AGN) contribution to its flux.  Its apparent (lensed) star-formation rate is 2417 \\Myr~(\\h0 = 75 km s$^{-1}$Mpc$^{-1}$, \\q0 = 0.1; with extinction correction) from the 1500\\AA~continuum (Pettini et al. 2000, hereafter P2000).  Throughout this paper, unless otherwise noted, a cosmology of (\\h0 = 75 km s$^{-1}$Mpc$^{-1}$, \\q0 = 0.1, $\\Lambda =0$) is assumed. ",
        "conclusions": "While it is difficult to draw general conclusions from observations of one object, the rest-frame optical spectrum of MS1512-cB58 affords us a preview of near-IR spectroscopy of LBGs.  As expected (Seitz et al. 1998, P2000), cB58 does not appear to be a ``proto-galaxy'', but a galaxy with significant metals produced in at least two phases.  We find that the properties of MS1512-cB58 are consistent with the interpretation of LBGs as progenitors of modern-day elliptical galaxies.  Previous studies have also shown that the mass, SFR, and extinction of LBGs are appropriate in that context (i.e. P98).  We have obtained the first measure of metal abundance in an LBG from the promising $R_{23}$~method.  The metallicity of $1/3$~solar is in the range expected at this redshift from hierarchical models of elliptical galaxy evolution (Thomas 1999).  Ellipticals produce most of their stars, and hence metals, in the first $\\sim 0.5$~or 1 Gyr (Rocca-Volmerange \\& Fioc, 2000), and so at $z\\sim 3$~they do not resemble low redshift, extremely metal poor galaxies, but rather older low-z starbursts.  Future LBG surveys with NIRSPEC will explore this connection (Pettini et al. 2000b).  The ease with which our spectra were obtained (1.3 hours of telescope time) demonstrates the great potential for observation of typical ($K\\simgt 20$) LBGs with NIRSPEC or similar instruments."
    },
    "0002/astro-ph0002022_arXiv.txt": {
        "abstract": "\\noindent We present an analysis of $\\sim$390 ksec of data of the Z source GX~340+0 taken during 24 observations with the {\\em Rossi\\,X\\,-ray\\,Timing\\,Explorer} satellite. We report the discovery of a new broad component in the power spectra. The frequency of this component varied between 9 and 14 Hz, and remained close to half that of the horizontal branch quasi-periodic oscillations (HBO). Its rms amplitude was consistent with being constant around $\\sim$5\\%, while its FWHM increased with frequency from 7 to 18 Hz. If this sub-HBO component is the fundamental frequency, then the HBO and its second harmonic are the second and fourth harmonic component, while the third harmonic was not detected. This is similar to what was recently found for the black hole candidate XTE~J1550--564. The profiles of both the horizontal- and the normal branch quasi-periodic oscillation peaks were asymmetric when they were strongest. We describe this in terms of a shoulder component at the high frequency side of the quasi-periodic oscillation peak, whose rms amplitudes were approximately constant at $\\sim$4\\% and $\\sim$3\\%, respectively. The peak separation between the twin kHz quasi-periodic oscillations was consistent with being constant at 339$\\pm$8 Hz but a trend similar to that seen in, e.g. Sco~X--1 could not be excluded. We discuss our results within the framework of the various models which have been proposed for the kHz QPOs and low frequency peaks. ",
        "introduction": "\\label{intro} \\noindent GX\\,340+0 is a bright low-mass X-ray binary (LMXB) and a Z source (Hasinger \\& van der Klis 1989).  The Z-shaped track traced out by Z sources in the X-ray color-color diagram or hardness-intensity diagram (HID) is divided into three branches: the horizontal branch (HB), the normal branch (NB), and the flaring branch (FB). The power spectral properties and the HID of GX\\,340+0 were previously described by van Paradijs et al. (1988) and Kuulkers \\& van der Klis (1996) using data obtained with the EXOSAT satellite, by Penninx et al. (1991) using data obtained with the Ginga satellite, and by Jonker et al. (1998) using data obtained with the {\\em Rossi\\,X\\,-ray\\,Timing\\,Explorer} {\\em(RXTE)} satellite. An extra branch trailing the FB in the HID has been described by Penninx et al. (1991) and Jonker et al. (1998). When the source is on the HB or on the upper part of the NB, quasi-periodic oscillations (QPOs) occur with frequencies varying from 20--50 Hz: the horizontal branch quasi-periodic oscillations or HBOs (Penninx et al. 1991; Kuulkers \\& van der Klis 1996; Jonker et al. 1998). Second harmonics of these HBOs were detected by Kuulkers \\& van der Klis (1996) and Jonker et al. (1998) in the frequency range 73--76 Hz and 38--69 Hz, respectively. In the middle of the NB, van Paradijs et al. (1988) found normal branch oscillations (NBOs) with a frequency of 5.6 Hz.  Recently, Jonker et al. (1998) discovered twin kHz QPOs in GX\\,340+0. These QPOs have now been seen in all six originally identified Z sources (Sco~X--1, van der Klis et al. 1996; Cyg X--2, Wijnands et al 1998a; GX~17+2, Wijnands et al. 1997b; GX~349+2, Zhang et al. 1998; GX~340+0, Jonker et al. 1998; GX~5--1, Wijnands et al. 1998b; see van der Klis 1997, 1999 for reviews), but not in Cir~X--1, which combines Z source and atoll source characteristics (Oosterbroek et al. 1995; Shirey, Bradt, and Levine 1999; see also Psaltis, Belloni, \\& van der Klis 1999).  \\par \\noindent In the other class of LMXBs, the atoll sources (Hasinger \\& van der Klis 1989), kHz QPOs are observed as well (see van der Klis 1997, 1999 for reviews).  Recently, also HBO-like features have been identified in a number of atoll sources (4U\\,1728--34, Strohmayer et al. 1996, Ford \\& van der Klis 1998, Di Salvo et al. 1999; GX13+1, Homan et al. 1998; 4U~1735--44, Wijnands et al. 1998c; 4U~1705--44, Ford, van der Klis, \\& Kaaret 1998; 4U\\,1915--05, Boirin et al. 1999; 4U~0614+09, van Straaten et al. 1999; see Psaltis, Belloni, \\& van der Klis 1999 for a summary). Furthermore, at the highest inferred mass accretion rates, QPOs with frequencies near 6 Hz have been discovered in the atoll sources 4U\\,1820--30 (Wijnands, van der Klis, \\& Rijkhorst 1999c), and XTE~J1806--246 (Wijnands \\& van der Klis 1998e, 1999b; Revnivtsev, Borozdin, \\& Emelyanov 1999), which might have a similar origin as the Z source NBOs.  \\par \\noindent At low mass accretion rates the power spectra of black hole candidates, atoll, and Z sources show similar characteristics (van der Klis 1994a,b).  Wijnands \\& van der Klis (1999a) found that the break frequency of the broken power law which describes the broad-band power spectrum, correlates well with the frequency of peaked noise components (and sometimes narrow QPO peaks) observed in atoll sources (including the millisecond X-ray pulsar SAX\\,J1808.4--3658; Wijnands \\& van der Klis 1998d, Chakrabarty \\& Morgan 1998), and black hole candidates.  The Z sources followed a slightly different correlation. In a similar analysis, Psaltis, Belloni, \\& van der Klis (1999) have pointed out correlations between the frequencies of some of these QPOs and other noise components in atoll sources, Z sources, and black hole candidates, which suggests these phenomena may be closely related across these various source types, or at least depend on a third phenomenon in the same manner. Because of these correlations, models describing the kHz QPOs which also predict QPOs or noise components in the low-frequency part of the power spectrum can be tested by investigating this low-frequency part.  \\par \\noindent In this paper, we study the full power spectral range of the bright LMXB and Z source GX~340+0 in order to further investigate the similarities between the atoll sources and the Z sources, and to help constrain models concerning the formation of the different QPOs. We report on the discovery of two new components in the power spectra of GX~340+0 with frequencies less than 40 Hz when the source is on the left part of the HB. We also discuss the properties of the NBO, and those of the kHz QPOs. ",
        "conclusions": "\\noindent In the present work we combined all {\\em RXTE} data presently available for the Z source GX~340+0 using our new selection method based on the frequency of the HBO peak. This allowed us to distinguish two new components in the low-frequency part of the power spectrum. \\par \\noindent These two extra components were strongest when the source was at the lowest count rates on the HB (see Fig.~\\ref{fig_HIDs}), between $\\rm{S_z}=$ 0.48--0.73, i.e., at the lowest inferred $\\dot{M}$.  The frequency of one of these components, the sub-HBO component, is close to half the frequency of the HBO component. The frequency ratio was consistent with being constant when the frequency of the sub-HBO changed from 9 to 14 Hz. A similar feature at sub-HBO frequencies has been reported by van der Klis et al. (1997) in Sco\\,X--1. Since the frequency of this component is close to twice the predicted Lense-Thirring (LT) precession frequency for rapidly rotating neutron stars (Stella \\& Vietri 1998), we shall discuss the properties of this component within this framework.  \\par \\noindent The other component we discovered, the HBO shoulder component, was used to describe the strong excess in power in the HBO profile towards higher frequencies.  If this shoulder component is related to the HBO and not to a completely different mechanism which by chance results in frequencies close to the frequency of the HBO, it can be used to constrain the formation models of the HBO peak. We demonstrated that both the HBO and the NBO have a similar asymmetric profile. In the NBO this was previously noted by Priedhorsky et al. (1986) in Sco~X--1. We shall consider the hypothesis that the formation of this shoulder is a common feature of the two different QPO phenomena, even if the two peaks themselves perhaps occur due to completely different physical reasons. \\par \\noindent Our results on the kHz QPOs based on more extensive data sets at three different epochs and using the new HBO selection method are consistent with those of Jonker et al. (1998). We discuss the properties of the kHz QPOs within the framework of precessing Keplerian flows (Stella \\& Vietri 1999), the sonic point model (Miller, Lamb, \\& Psaltis 1998), and the transition layer model described by Osherovich \\& Titarchuk (1999), and Titarchuk, Osherovich, \\& Kuznetsov (1999).  \\subsection{Comparison with other observations} \\noindent In various LMXBs, QPOs have been found whose profiles are clearly not symmetric. Belloni et al. (1997) showed that for the black hole candidate (BHC) GS~1124--68 the QPO profiles are asymmetric, with a high frequency shoulder. Dieters et al. (1998) reported that the 2.67 Hz QPO of the BHC 4U~1630--47 was also asymmetric with a high frequency shoulder. In the Z source Sco~X--1 the NBO profile was also found to be asymmetric (Priedhorsky et al. 1986). It is clear that asymmetric shapes of the QPO profiles are frequently observed in LMXBs and are not restricted to either the black hole candidates or the neutron star systems.\\par \\noindent In the BHCs GS~1124--68 (Belloni et al. 1997) and XTE~J1550--564 (Homan et al. 1999) several QPOs were discovered which seem to be harmonically related in the same way as we report for GX~340+0, i.e. the third harmonic is not detected, while the first, the second and the fourth harmonic are. If this implies that these QPOs are the same, models involving the magnetic field of the neutron star for their origin could be ruled out. The time lag properties of the harmonic components of the QPOs in XTE~J1550--564 are complex and quite distinctive (Wijnands, Homan, \\& van der Klis 1999). In GX~340+0 no time lags of the harmonic components could be measured, but the time lags measured in the HBO in the similar Z source GX~5--1 (Vaughan et al. 1994) are quite different. \\par \\noindent In order to study in more detail the relationship found by Wijnands \\& van der Klis (1999) between the QPOs and the noise break frequency in the power spectrum of LMXBs, we fitted the LFN component using a broken power law. To determine the value for the break frequency we fixed the parameters of all other components to the values found when using a cut-off power law to describe the LFN. Wijnands \\& van der Klis (1999a) reported that the Z sources did not fall on the relation between the break and QPO frequency established by atoll sources and black hole candidates. They suggested that the Z source LFN is not similar to the atoll HFN but the noise component found in Sco~X--1 at sub-HBO frequencies is. By using the centroid frequency of that peaked noise component as the break frequency instead of the LFN break frequency, the HBO frequencies did fall on the reported relation. On the other hand, we find that using the sub-HBO frequency instead of the HBO frequency together with the LFN break frequency, the Z source GX~340+0 also falls exactly on the relation.  Therefore, the suggestion made by Wijnands \\& van der Klis (1999a) that the strong band-limited noise in atoll and Z sources have a different origin is only one of the two possible solutions to the observed discrepancy. Our proposed alternative solution is that the Z and atoll noise components are the same, but that it is the sub-HBO in Z sources which corresponds to the QPO in atoll sources. An argument in favour of the noise components in Z and atoll sources being the same is that the cut-off frequency of the LFN component increased as a function of $\\rm{S_z}$, in a similar fashion as the frequency associated with the atoll high frequency noise (van der Klis 1995, Ford \\& van der Klis 1998, van Straaten et al. 1999). \\par \\noindent Following Psaltis, Belloni, \\& van der Klis (1999) we plotted the sub-HBO frequency against the frequency of the lower-frequency kHz QPO. The sub-HBO does not fall on the relation found by Psaltis, Belloni, \\& van der Klis (1999) between the frequency of the HBO and the lower-frequency kHz QPO frequency. Instead the data points fall between the two branches defined by the HBO-like QPO frequencies vs. the lower kHz QPO frequency at high frequencies (see Psaltis, Belloni, \\& van der Klis 1999).  \\subsection{HBO -- kHz QPO relations} \\subsubsection{Lense-Thirring precession frequency} \\noindent Stella \\& Vietri (1998) recently considered the possibility that the HBO is formed due to the LT precession of the nodal points of sligthly tilted orbits in the inner accretion disk, but as they already mentioned the Z sources GX~5--1 and GX~17+2 did not seem to fit in this scheme. For reasonable values of I/M, the neutron star moment of inertia divided by its mass, the observed frequencies were larger by a factor of $\\sim$2 than the predicted ones. Jonker et al. (1998) showed that for GX~340+0 the predicted frequency is too small by a factor of 3, if one assumes that the higher frequency peak of the kHz QPOs reflects the Keplerian frequency of matter in orbit around the neutron star, and that the mean peak separation reflects the neutron star spin frequency. Using the same assumptions Psaltis et al. (1999) also concluded that a simple LT precession frequency model is unable to explain the formation of HBOs in Z sources.  \\par \\noindent Detailed calculations of Morsink \\& Stella (1999) even worsen the situation, since their calculations lower the predicted LT frequencies. They find that the LT precession frequencies are approximately a factor of two too low to explain the noise components at frequencies $\\sim$20--35 Hz observed in atoll sources (4U~1735--44, Wijnands \\& van der Klis 1998c; 4U~1728--34, Strohmayer et al 1996, Ford \\& van der Klis 1998). Stella \\& Vietri (1998) already put forward the suggestion that a modulation can be produced at twice the LT precession frequency if the modulation is produced by the two points where the inclined orbit intersects the disk plane (although they initially used this for explaining the discrepancy of a factor of two between the predicted and the observed LT precession frequencies for the Z sources).  \\par \\noindent The sub-HBO peaked noise component we discovered could be harmonically related to the HBO component. If the sub-HBO is the second harmonic of the fundamental LT precession frequency, as needed to explain the frequencies in the framework of the LT precession model where the neutron star spin frequency is approximately equal to the frequency of the kHz QPO peak separation, the HBO must be the fourth and the harmonic of the HBO must be the eighth harmonic component, whereas the sixth and uneven harmonics must be much weaker. This poses strong (geometrical) constraints on the LT precession process. On the other hand, if the HBO frequency is twice the LT precession frequency, which implies a neutron star spin frequency of $\\sim$900 Hz (see Morsink \\& Stella 1999), the frequency of the sub-HBO component is the LT precession frequency, and the frequency of the second harmonic of the HBO is four times the LT precession frequency. In that case only even harmonics and the LT precession frequency are observed.  \\subsubsection{Magnetospheric beat frequency and radial-flow models} \\label{random} \\noindent In this section, we discuss our findings concerning the QPOs and the LFN component in terms of the magnetic beat frequency model where the QPOs are described by harmonic series (e.g. Shibazaki \\& Lamb 1987). \\par \\noindent If the sub-HBO frequency is proven not to be harmonically related to the HBO, the sub-HBO peak might be explained as an effect of fluctuations entering the magnetospheric boundary layer periodically. Such an effect will be strongest at low HBO frequencies since its power density will be proportional to the power density of the LFN (Shibazaki \\& Lamb 1987). If it is the fundamental frequency and the HBO its first overtone then the magnetospheric beat frequency model proposed to explain the HBO formation (Alpar \\& Shaham 1985; Lamb et al. 1985) is not strongly constrained.  \\par \\noindent Within the beat frequency model the high frequency shoulder of the HBO peak can be explained as a sign of radial drift of the blobs as they spiral in after crossing the magnetospheric boundary layer (Shibazaki \\& Lamb 1987). Shibazaki \\& Lamb (1987) describe another mechanism which may produce a high frequency shoulder. Interference between the LFN and the QPO caused by a non uniform phase distribution of the blobs will also cause the QPO to become asymmetric. This effect will be strongest when the LFN and the QPO components overlap, as observed. Finally, an asymmetric initial distribution of frequencies of the blobs when entering the magnetospheric boundary layer may also form an asymmetric HBO peak.  \\par \\noindent The changes in the power law index of the LFN as a function of photon energy can be explained by varying the width or the steepness of the lifetime distribution of the blobs entering the magnetic boundary layer (Shibazaki \\& Lamb 1987). The decrease in increase of both the fractional and absolute rms amplitude of the HBO as a function of energy towards higher frequencies (Fig.~\\ref{rmsratio}) also constrains the detailed physical interactions occurring in the boundary layer.  \\par \\noindent Fortner et al. (1989) proposed that the NBO is caused by oscillations in the electron scattering optical depth at the critical Eddington mass accretion rate. How a high frequency shoulder can be produced within this model is not clear. Both the HBO and the NBO shoulder components were detected when the rms amplitude of the HBO and the NBO was highest. In case of the NBO, this may be a result of the higher signal to noise. Since the rms amplitude of the NBO shoulder component is consistent with being $\\sim$2/3 of the NBO rms amplitude (see Table~\\ref{nbo_tab}), combining more observations should increase the range over which this shoulder component is detected, if this ratio is constant along $\\rm{S_z}$. In case of the HBO the two components seem to merge. While the fractional rms amplitude of the HBO shoulder component increased that of the HBO decreased. When the fractional rms amplitudes were comparable, the HBO was fitted with one Lorentzian. The rms amplitude of both shoulder components increased in a similar way as the rms amplitudes of the NBO and the HBO with photon energy. So, the formation of these shoulder components seems a common feature of both QPO forming mechanisms.  \\subsubsection{Radial oscillations in a viscous layer} In Sco~X--1, Titarchuk, Osherovich, \\& Kuznetsov (1999) interpreted the extra noise component in the power spectra (van der Klis et al. 1997) as due to radial oscillations in a viscous boundary layer (Titarchuk \\& Osherovich 1999). If the noise component in Sco~X--1 is the sub-HBO component in GX~340+0, the model of Titarchuk \\& Osherovich (1999) can be applied to the frequencies and dependencies we found for the sub-HBO component in GX~340+0. Fitting our data to the relation between the frequency of the extra noise component and the Keplerian frequency, using the parameters and parametrization given by Titarchuk, Osherovich, \\& Kuznetsov (1999), we obtained a value of $C_N=15$ for GX~340+0. This value is much larger than the value obtained for Sco~X--1 (9.76). According to Titarchuk \\& Osherovich (1999) a higher $C_N$ value implies a higher viscosity for the same Reynold's number.  \\subsection{KHz QPOs and their peak separation} \\noindent Recently, Stella \\& Vietri (1999) have put forward a model in which the formation of the lower kHz QPO is due to the relativistic periastron precession (apsidal motion) of matter in (near) Keplerian orbits. The frequency of the upper kHz QPO peaks is the Keplerian frequency of this material. The peak separation is then equal to the radial frequency of matter in a nearly circular Keplerian orbit, and is predicted to decrease as the Keplerian frequency increases and approaches the predicted frequency at the marginally stable circular orbit. This model can explain the decrease in peak separation as observed in various sources (see Section~\\ref{khz_res}).  \\par \\noindent Beat frequency models stating that the upper kHz QPO peak is formed by Keplerian motion at a preferred radius in the disk (e.g. the sonic point radius, Miller, Lamb, \\& Psaltis 1998), whereas the lower kHz QPO peak formed at the frequency of the beat between the neutron star spin frequency and this Keplerian frequency, cannot in their original form explain the decrease in peak separation in these two sources. A relatively small extension of the model (Lamb, Miller, \\& Psaltis 1998) can, however, produce the observed decrease in peak separation.\\par \\noindent Osherovich \\& Titarchuk (1999) developed a model in which the kHz QPOs arise due to radial oscillations of blobs of accreting material at the magnetospheric boundary. The lower kHz QPO frequency is in their model identified with the Keplerian frequency. Besides this QPO two eigenmodes are identified whose frequencies coincide with the upper kHz QPO peak frequency and the frequency of the HBO component in the power spectra of Sco~X--1 (Titarchuk \\& Osherovich 1999). Interpreting our findings within this framework did not result in stringent constraints on the model.  \\par \\noindent We found that the peak separation is consistent with being constant (Fig.~\\ref{kHz_vs_HBO_two_selection} A and B), but neither a decrease towards higher $\\dot{M}$ as in Sco~X--1, 4U~1608--52, 4U~1735--44, 4U~1702--429, and 4U~1728--34 nor a decrease towards lower $\\dot{M}$, as predicted by Stella \\& Vietri (1999) can be ruled out. If the model of Stella \\& Vietri turns out to be the right one the mass of the neutron star most likely is in the range of 1.8 to 2.2 $M_{\\odot}$ (see Fig.~\\ref{kHz_vs_HBO_two_selection} B). This is in agreement with the mass of Cyg~X--2 derived by Orosz \\& Kuulkers (1999), and with the masses of the neutron stars derived when interpreting the highest observed kHz QPO frequencies as due to motion at or near the marginally stable orbit (Kaaret, Ford, \\& Chen 1997; Zhang, Strohmayer, \\& Swank 1997)."
    },
    "0002/astro-ph0002079_arXiv.txt": {
        "abstract": "We have used infrared polarimetric imaging with NICMOS to determine precisely the position of the star that illuminates (and presumably generated) the bipolar, pre-planetary reflection nebula RAFGL 2688 (the Egg Nebula). The polarimetric data pinpoint the illuminating star, which is not detected directly at wavelengths $\\le$ 2 $\\mu$m, at a position well within the dark lane that bisects the nebula, 0\\farcs55 ($\\sim550$ AU) southwest of the infrared peak which was previously detected at the southern tip of the northern polar lobe. The inferred position of the central star corresponds to the geometric center of the tips of the four principle lobes of near-infrared H$_2$ emission; identifying the central star at this position also reveals the strong point symmetric structure of the nebula, as seen both in the intensity and polarization structure of the polar lobes. The polarimetric and imaging data indicate that the infrared peak directly detected in the NICMOS images is a self-luminous source and, therefore, is most likely a distant binary companion to the illuminating star. Although present theory predicts that bipolar structure in pre-planetary and planetary nebulae is a consequence of binary star evolution, the separation between the components of the RAFGL 2688 binary system, as deduced from these observations, is much too large for the presence of the infrared companion to have influenced the structure of the RAFGL 2688 nebula. ",
        "introduction": "The bipolar structures exhibited by a substantial fraction of the known planetary nebulae likely arise during the last, rapid, pre-planetary nebula (PPN) stage of evolution of intermediate-mass (1--8 M$_\\odot$) stars off the asymptotic giant branch (AGB). A popular, albeit largely untested, model for such bipolarity is that the central AGB star possesses a companion that aids in the buildup of a dense, dusty equatorial torus surrounding the central star (e.g., \\cite{soke1998}).  Alternatively, the fossil remnant of a $\\beta$ Pic-like main-sequence disk may bear responsibility for triggering bipolarity during post-main sequence evolution (\\cite{kast1995}). Whatever the mechanism that abets their formation, bipolar PPN typically show two bright reflection lobes separated by a dark dust lane.  The star that illuminates the polar lobes presumably is located at or near the center of the equatorial, dust torus. While this geometry obscures the central star along our direct line of sight, photons readily escape the nebular core in the polar directions and subsequently are scattered by dust grains located primarily in the walls of the rarefied, expanding lobes.  As even the lobe walls tend to be optically thin in the near-infrared, such photons can be singly scattered out of the nebula into our line of sight.  Single scattering produces polarized light that contains a record of the original direction of the unpolarized light source; therefore, polarimetric maps of such polarized nebulae contain clues as to the locations of their illuminating sources, even if those stars lie hidden inside dust lanes.  Recent direct imaging of RAFGL 2688 (the Egg Nebula) with the Near Infrared Camera and Multi-Object Spectrometer (NICMOS) aboard the Hubble Space Telescope ({\\it HST}) (\\cite{saha1998}) revealed a compact red source just south of the bottom of the northern reflection lobe.  However, initial analysis of the polarimetric maps from NICMOS indicated that this red source was not the primary illuminator of the reflection nebulosity; this object is most likely a companion to the post-AGB star that lurks in the core of the Egg Nebula.  From a preliminary examination of the 2.0 $\\mu$m polarimetric map, Sahai et al. suggested that the obscured, post-AGB star was located $\\simeq$ 750 AU (0\\farcs75) south of the red companion.  In this paper, we present a rigorous analysis of the 2.0 $\\mu$m polarization map of RAFGL 2688 obtained by NICMOS.  We determine the precise position of the post-AGB star in the core, assess the relationship of the red source to the illuminator star, and discuss the implications of this work for understanding the formation of the Egg Nebula and of other bipolar PPN. ",
        "conclusions": "From a detailed analysis of the polarimetric images obtained using NICMOS and the {\\it HST}, we have precisely determined the position of the post-AGB star in the waist of the Egg Nebula and the projected orientation of the polar axis (PA 12$^\\circ$) of this bipolar system. This post-AGB star, which illuminates the Egg Nebula, falls point-symmetrically at the center of the molecular hydrogen emission regions that mark the waist and the polar lobes of the nebula. We find that this star lies 550 AU in projected distance, and perhaps 1000 AU in physical distance, from the star previously identified (\\cite{saha1998}) at the southern tip of the northern polar lobe. Thus, these data provide clear evidence for the presence of an optically obscured, widely spaced binary system near the core of the bipolar, pre-planetary nebula RAFGL 2688.  However, the separation between these components is orders of magnitude larger than required by models postulating that companions to AGB stars trigger the production of bipolar planetary nebulae."
    },
    "0002/astro-ph0002415_arXiv.txt": {
        "abstract": "Based on the analysis of a large sample of RXTE/PCA observations of several black hole binaries in the low spectral state we show that a correlation exists between the spectral parameters and characteristic noise frequency. In particular, the amplitude of reflection increases and the slope of Comptonized radiation steepens as the noise frequency increases. We consider possible implications of these results on the accretion flow models and discuss a possible observational test aimed to discriminate between different geometries of the accretion flow. ",
        "introduction": "\\subsubsection{Observations and data reduction}  The results presented below are based on the publicly available data of Cyg X--1, GX339--4 and GS1354--644 observations with the Proportional Counter Array aboard the Rossi X-ray Timing Explorer (\\cite{rxte}) from 1996--1998 during the low (hard) spectral state of the sources. In total our sample included $\\approx 60$ observations of these sources. The energy and power density spectra were averaged for each individual observation. The energy spectra were extracted from the ``Standard Mode 2'' data and ARF and RMF were constructed using standard RXTE FTOOLS v.4.2 tasks. The ``VLE'' and ``Q6'' models were used for the background calculation with the preference being given to the ``VLE'' model when possible. A uniform systematic error of 0.5\\% was added quadratically to the statistical error in each energy channel. The power spectra were calculated in the 2--16 keV energy band and in the $\\approx 0.005-32$ Hz frequency range  following the standard X--ray timing technique and normalized to units of squared fractional rms.        \\subsubsection{The spectral model}   The energy spectra were fit with a model consisting of a power law with a superimposed reflected continuum (pexrav model in XSPEC, \\cite{pexrav}) and an intrinsically narrow line at $\\sim 6.4-6.7$ keV. The centroid energy of the line  was a free parameter of the fit. For all sources the binary system inclination was fixed at $i=50$ degrees, the iron abundance -- at the solar value of $A_{\\rm Fe}=3.3\\cdot 10^{-5}$ and the low energy  absorption --  at $N_{\\rm H}=6\\cdot 10^{21}$ cm$^{-2}$.  The effects of ionization were not included. In order to approximately include in the model the smearing of the reflection features due to motion in the accretion flow the reflection continuum and line were convolved with a Gaussian, which width was a free parameter of the fit. The spectra were fitted in the 4--20 keV energy range. This spectral model is identical to that applied by \\cite{paper1} to  the Cyg X-1 data, except that in the present analysis we let the line centroid energy be a free parameter of the fit. Strength of the reflected component is characterized in the pexrav model by the reflection scaling factor $R$, which approximately measures the solid angle subtended by the reflecting media as seen from the source of the primary radiation, $R\\approx \\Omega_{refl}/2\\pi$, so that $R=1$ for an isotropic point source above an infinite optically thick slab.   The spectral model is obviously oversimplified. Moreover, several important parameters, such as the binary system inclination angle and the iron abundance were fixed at fiducial values. Therefore the best fit values  do not necessarily represent the exact values of the physically interesting parameters.  Particularly subject to the uncertainties due to the choice of the spectral model is the reflection scaling factor $R\\sim\\Omega/2\\pi$. This might explain the values of $R$ exceeding unity obtained for some of the spectra. However, as was shown by \\cite{paper1} the model correctly ranks the spectra according to the strength of the reflected component and the slope of the underlying power law.    \\begin{figure}[t] \\centerline{ \\epsfxsize 9. cm \\epsffile{refl_slope_mod_ab.ps} } \\caption{The best fit values of the photon index of the underlying power law plotted vs. reflection scaling factor for Cyg X-1, GX339--4 and GS1354--644. The solid line shows the dependence $\\Gamma(R)$ of the photon index of the Comptonized radiation $\\Gamma$ on the strength of the reflected emission $R$ expected in the disk--spheroid model assuming the disk albedo $a=0.1$, Thomson optical depth of the cloud $\\tau_T=1$ and the ratio of the temperature of the seed photons to the electron temperature $T_{bb}/T_e=10^{-4}$. The dashed line shows $\\Gamma(R)$ dependence expected in the plasma ejection model (\\cite{belob}) for $a=0.15$, $\\tau_T=1$, $T_{bb}/T_e=3\\cdot 10^{-3}$ and $\\mu_s=0.3$. \\label{refl_slope} } \\end{figure}   \\begin{figure}[t] \\centerline{ \\epsfxsize 9.0 cm \\epsffile{qpo_refl.ps} } \\caption{The QPO centroid frequency  plotted vs. reflection scaling factor for the same set of observations of Cyg X-1, GX339--4 and GS1354--644 as in Fig.\\ref{refl_slope} and \\ref{refl_gsmo}. \\label{refl_qpo} } \\end{figure}      \\subsubsection{Power spectra approximation}   The power density spectra in units of $rms^2/Hz$  were fit in the 0.005--15 Hz frequency range with a model, consisting of two broken constant components and a Lorentzian profile $P_\\nu=\\frac{P_0}{1+(2(\\nu-\\nu_0)/\\Delta\\nu)^2}$. We use centroid frequency of the  Lorentzian profile  as a characteristic noise frequency. Its dependence on the break frequency of the band limited noise (the first broken constant component in the model) is shown in Fig.\\ref{pds}.       \\subsubsection{Correlation between the spectral parameters and the noise frequency} The results of the energy and power spectra approximation are presented in Fig. \\ref{refl_slope} and \\ref{refl_qpo} showing dependence of the slope of the primary emission and characteristic noise frequency on the amplitude of the reflected component. The noise frequency and the spectral parameters are clearly correlated. An increase of the amplitude of the reflected component is accompanied with a steepening of the spectrum of the primary emission and an increase of the noise frequency. ",
        "conclusions": "In the low spectral state of black hole candidates the amplitude of reflection $R$, slope of the Comptonized radiation $\\Gamma$ and characteristic noise frequency change in a correlated way. The correlation between the slope and reflection indicates that the accretion disk is the primary source of the soft seed photons for Comptonization. We considered two idealized models, the disk--spheroid and plasma ejection model, and showed that with appropriate tuning of the parameters either model can explain the observed  $\\Gamma(R)$ dependence. The RXTE data also hints at a correlated change of the width of relativistic smearing of the reflection features. If confirmed by observations with better energy resolution, this correlation may help to distinguish between different geometries of the accretion flow."
    },
    "0002/astro-ph0002309_arXiv.txt": {
        "abstract": "We  present results of  extensive model calculations  of disk galaxy evolution within an  hierarchical inside-out formation scenario.  We first   compare   properties of  the  dark   halos   identified in a cosmological N-body simulation  with predictions of a  seminumerical method   based on  an extended    collapse model  and  find a   good agreement.  We then describe detailed modelling of the formation and evolution of disks within   these growing halos and  predictions for the   main properties,   correlations and evolutionary   features of normal  disk   galaxies.  The  shortcomings   of   the scenario  are discussed. ",
        "introduction": "The hierarchical   cosmic structure  formation picture based    on the inflationary  cold  dark matter  (CDM) provides  a solid framework for models  of galaxy  formation and evolution.   On  the other  hand, the unprecedented observations of  galaxies at different redshifts make it possible to probe and constrain these models.  Here we discuss some of the  results obtained with  a self-consistent scenario  of disk galaxy formation and evolution within the context of the hierarchical picture (the extended collapse scenario). ",
        "conclusions": ""
    },
    "0002/astro-ph0002145_arXiv.txt": {
        "abstract": "We develop a numerical chemical model allowing for radial flows of gas, with the aim to analyse the possible role of gas flows in the chemical evolution of the Galactic Disc. The dynamical effects of the Galactic Bar on the radial gas profile of the Disc are especially addressed. ",
        "introduction": "The observed chemical and spectro--photometric properties of galaxies are one of the main sources of information for our understanding of galaxy formation and evolution. The corresponding theoretical modelling involves star formation (SF) as a basic ingredient. Unfortunately, this process is rather poorly known on the large scales relevant to galaxy evolution. Portinari \\& Chiosi (1999, hereinafter PC99) analysed the effects of adopting different SF laws in a chemical model for the Galactic Disc, a system which we can study in great detail. In this paper we address another phenomenon which can bear interesting effects on the chemical evolution of galaxies: radial gas flows. A few papers in literature (Section~\\ref{previous}) demonstrate that radial gas flows influence chemical models for the Disc, especially in their predictions on the metallicity gradient. It is therefore interesting to discuss the radial profile of the Disc with models including also radial flows, in addition to various options for the SF law. In particular, radial flows can help to overcome some difficulties that ``static'' models find in reproducing, at the same time, the metallicity gradient and the radial gas profile of the Disc (PC99).  We develop a chemical model with radial gas flows as a multi--dimensional extension of the model of Portinari \\etal (1998, hereinafter PCB98). The model is described in Section~\\ref{discrete} and in the appendices. In Section~\\ref{radfloweffects} we discuss the general qualitative effects of superposing radial flows upon a chemical model. In Section~\\ref{bestfit} we present models for the Galactic Disc with radial gas flows and different SF laws, showing that radial flows provide an alternative or additional dynamical effect to the ``inside--out'' formation scenario to explain the metallicity gradient. Section~\\ref{bar} is dedicated to qualitative simulations of the dynamical effects of the Galactic Bar upon the gas distribution, with the aim to reproduce the molecular ring around 4 kpc, which static models cannot account for (PC99). Section~\\ref{conclusions} contains a final summary and conclusions. ",
        "conclusions": "\\label{conclusions}  From the results of static chemical models, PC99 underlined the need to introduce radial flows to explain some features of the Galactic Disc. In fact, static models are unable to reproduce, at the same time, both the metallicity gradient and the radial gas profile; in particular, the peak corresponding to the molecular ring at 4--6~kpc is likely to be a consequence of gas drifts induced by the dynamical influence of the Galactic Bar. Therefore, in the present paper we introduced a new chemical model including radial gas flows, developed as a multi--dimensional generalization of the original static model (\\S\\ref{discrete}). Our model is conceived so as to adapt to any imposed radial velocity profile, describing both inflows and outflows in any part of the disc. The model is carefully tested against instability problems and spatial resolution, by comparing it to suitable exact analytical cases (Appendix~B). In this paper we applied the model to the Galactic Disc; more in general, such models allowing for gas drifts are meant to be used as fast and handy interfaces between detailed dynamical galaxy models (predicting the velocity profiles) and parametric chemical and spectro--photometric models.  An overview of the behaviour of chemical models with radial inflows of gas shows that these provide an alternative ``dynamical'' assumption to the inside--out disc formation scenario to explain the metallicity gradient (\\S\\ref{radfloweffects}). With radial gas flows, the model can reproduce the metallicity gradient even in the case of a Schmidt or an {\\mbox{Oort--type}} SF law, which were excluded in the case of static models (see PC99). In addition, it appears that even low radial flow velocities, well within observational limits and theoretical expectations (see~\\S\\ref{previous}), have non--negligible effects upon model predictions on the metallicity gradient and moreover on the gas distribution. In particular, if radial gas inflows are allowed for, a metallicity gradient can coexist with a high gas fraction in the inner regions, at odds with simple static models. This is indispensable to reproduce the observed gas distribution in the inner Galaxy (see point 2 below). The remarkable effects of even slow radial flows upon observable quantities, mainly upon the gas distribution, should be kept in mind as a caveat when comparing real galaxies to simple analytical models which predict a one--to--one relationship between metallicity and gas fraction (e.g.\\ Tinsley 1980). Our models show that small dynamical effects, like slow gas flows, can easily make real systems depart from the behaviour of simple models.  With our model it is possible to mimic the dynamical influence of the Galactic Bar and reproduce the peak in the gas distribution around 4~kpc (\\S\\ref{bar}). Two scenarios, related to two different models for the structure of the Bar, are qualitatively suggested. \\begin{description} \\item[{\\bf A.}] With a Schmidt or an Oort--type SF law, slow radial inflows in the disc pile up gas inward down to {\\mbox{$r=3.5-4$~kpc}}. Here, the Bar CR radius is found and the gas is quickly swept inward from CR toward an ILR, which causes the drop in the gas profile at 3.5~kpc. \\item[{\\bf B.}] With a SF law like that by Dopita \\& Ryder (1994), smaller inflow rates suffice to reproduce the metallicity gradient, leading to a lower concentration of gas in the inner regions than in the previous case. The peak corresponding to the molecular ring can be reproduced with a Bar CR around 2.5~kpc and its OLR around 4.5~kpc, so that all the gas external to $r=2.5$~kpc tends to pile up around the OLR. \\end{description} Though these models are just qualitative and cannot describe the detailed dynamical process of Bar formation nor the evolution of the related gas flows to form the molecular ring, they provide two interesting indications. \\begin{enumerate} \\item Only when introducing the effects of the Bar, the model is able to reproduce the radial gas profile properly. The only possible exception resides in a particular combination of an Oort-type SF law with a radial inflow pattern whose velocity decreases inward (model {\\sf O10RFe}; this combination may lead to a peak of the gas distribution in the inner Galactic regions, closed to the observed molecular ring. But this particular, fortunate case does not diminish the general conclusions about the role of the Galactic Bar. \\item In any case (A or B above), overall radial inflows in the disc are indispensable to replenish the inner regions with enough gas that the observed molecular ring can form under the influence of the Bar. This seems to favour disc models with radial inflows, unless one assumes that the gas in the ring has some different origin (gas swept from the Bulge, or accreted later). \\end{enumerate} To investigate these issues any further, detailed gas--dynamical models are obviously required. Unfortunately, most studies on Bar--induced gas dynamics (see references in \\S\\ref{bar}) concentrate on the observed features of the very inner regions, such as the nuclear ring, the 3~kpc expanding arm, and so forth. Little discussion can be found about the effects of the Bar on more external regions, and on the formation of the molecular ring in particular: whether it is due to gas depletion inside CR as in our case A, or due to gas accumulation at some resonance (e.g.\\ Binney \\etal 1991, Fux 1999) as in our case B, or whether it just consists of two or more tightly wound spiral arms (e.g. Englmaier \\& Gerhard 1999). Further gas--dynamical studies suggesting detailed scenarios, time-scales, and velocity profiles for the formation of the molecular ring would be welcome, for the sake of including the effects of the Bar in chemical evolution models more consistently.  Further investigation of gas--dynamical models on the influence of the Bar on even larger scales (namely, outside its OLR) should be pursued as well, since this is a clue issue related to a claimed discrepancy between the characteristics of the Galactic Bar and the observed metallicity gradient. It is well known that barred galaxies display systematically shallower gradients than ordinary spirals (e.g.\\ Alloin \\etal 1981, Vila--Costas \\& Edmunds 1992, Martin \\& Roy 1994). This is likely a consequence of the radial mixing induced by bars; in fact, Martin \\& Roy (1994) found a correlation for external galaxies between the strength of a bar and the metallicity gradient. Taking this empirical relation at face value, the Galactic Bar with an axial ratio of {\\mbox{$\\sim 0.5$}} should induce a metallicity gradient of {\\mbox{--0.03~dex/kpc}}, much shallower than the observed one of --0.07~dex/kpc, which is typical of a {\\it normal} Sbc galaxy. To overcome such a puzzle, it has been suggested that the Galactic Bar must be very young ($<$1~Gyr), so that there was not enough time yet to flatten the gradient (Gummersbach \\etal 1998); but this is in conflict with other estimates of the Bar's age (e.g.\\ Sevenster 1997, 1999). Alternatively, we suggest that the discrepancy might be only apparent, since the Galactic Bar is quite small, and the Milky Way cannot be properly considered a barred spiral. It might be unlikely that the Bar can influence the metallicity gradient all over the Disc, as in really barred galaxies: Bar--induced radial drifts and corresponding chemical mixing are expected to occur from CR toward the ILR (inflows) and to the OLR (outflows; e.g.\\ Schwarz 1981, 1984; Friedli \\etal 1994). Present understanding of the Galactic Bar sets its OLR between 4.5 and 6~kpc (\\S\\ref{bar} and references therein), so in our models we presumed that the Bar induces negligible mixing beyond these radii, regardless of its age (see also Gerhard 1999). If the situation is as in the models we presented here, in fact, the metallicity gradient in the outer regions is unperturbed and just related to intrinsic Disc properties and/or large--scale viscous flows.  However, gas--dynamical simulations dedicated to the effects of the Galactic Bar over the whole Disc would be necessary, so as to investigate the relation between the Bar, radial mixing and the metallicity gradient more consistently. More in general, including the effects of {\\mbox{bar--induced}} radial flows in the picture of the chemical evolution of spiral galaxies might turn out to be of wide interest, since it is likely that all spirals develop at some point, or have developed in the past, some bar--like structure (Binney 1995). Infrared observations indeed reveal that a large fraction of spirals host a barred structure (e.g.\\ Eskridge \\etal 1999), and recent numerical simulations suggest that even weak bars or oval distortions may be able to induce radial drifts to form multiple gaseous rings at the corresponding Lindblad resonances (Jungwiert \\& Palou\\v{s} 1996). Bars could even drive secular evolution of spiral discs from late to early type (e.g.\\ Dutil \\& Roy 1999). Bar--driven radial gas flows might therefore play a fundamental role in the chemical evolution of spiral discs."
    },
    "0002/astro-ph0002373_arXiv.txt": {
        "abstract": "This review describes some general properites of proto-planetary nebulae with particular emphasis on the recent work of morpholgical studies.  The weight of observational evidence shows that proto-planetary nebulae (PPNe) are most certainly axisymmetric like planetary nebulae.  Recent work suggests two subclasses of PPNe optical morphology, DUst-Prominent Longitudinally-EXtended (DUPLEX) and Star-Obvious Low-level Elongated (SOLE). Radiative transfer models of an example DUPLEX PPN and SOLE PPN, presented here, support the interpretation that DUPLEX and SOLE are two physically distinct types of PPNe.  The DUPLEX PPNe and SOLE PPNe may well be the precursors to bipolar and elliptical PNe, respectively. ",
        "introduction": "The proto-planetary nebula (PPN;a.k.a. post-AGB or pre-PN) stage of evolution immediately precedes the planetary nebula (PN) stage.  The lifetime  for this phase is $\\; \\buildrel < \\over \\sim \\;$1000 years and marks the time from when the star was forced off the asymptotic giant branch (AGB) by intensive mass loss to when the central star becomes hot enough (T$_{\\rm eff} \\sim 3\\times 10^4$ K) to photoionize the neutral circumstellar shell (Kwok 1993). We refer the reader to Kwok (1993) for a comprehensive review of PPN. For short recent reviews, see Hrivnak (1997)  and van Winckel (1999). In this short review, I summarize the basic properties of PPN but focus primarily on the morphological studies because there have been numerous morphological studies in the past few years and because this particular conference is focused on morphology. ",
        "conclusions": "The observational evidence from a number of independent studies clearly shows that PPNe  are intrinsically  axisymmetric. Thus the axisymmetry that we observe in PNe must predate the PPNe stage.  Most likely the axisymmetry originates at the end of the AGB phase, because observations of the outer regions of AGB envelopes show a spherical symmetry.  With the variety of PNe morphologies discussed at this conference (e.g.  round, elliptical, bipolar and point symmetric), we must now  begin to ask: Do we see examples of PPNe with these corresponding subtle differences in morphologies? I think we are beginning to see these differences.  The DUPLEX and SOLE PPNe may well be the precursors for bipolar and elliptical PNe."
    },
    "0002/astro-ph0002003_arXiv.txt": {
        "abstract": "We describe a population of compact objects in the centre of the Fornax Cluster which were discovered as part of our 2dF Fornax Spectroscopic Survey. These objects have spectra typical of old stellar systems, but are unresolved on photographic sky survey plates. They have absolute magnitudes $-13<M_B<-11$, so they are 10 times more luminous than any Galactic globular clusters, but fainter than any known compact dwarf galaxies.  These objects are all within 30 arcminutes of the central galaxy of the cluster, NGC 1399, but are distributed over larger radii than the globular cluster system of that galaxy.  We suggest that these objects are either super-massive star clusters (intra-cluster globular clusters or tidally stripped nuclei of dwarf galaxies) or a new type of low-luminosity compact elliptical dwarf (``M32-type'') galaxy.  The best way to test these hypotheses will be to obtain high resolution imaging and high-dispersion spectroscopy to determine their structures and mass-to-light ratios. This will allow us to compare them to known compact objects and establish if they represent a new class of hitherto unknown stellar system. ",
        "introduction": "In cold dark matter (CDM) galaxy formation, small dense halos of dark matter collapse at high redshift and eventually merge to form the large virialised galaxy clusters seen today. The CDM model is very good at reproducing large-scale structure, but only very recently have the best numerical simulations (Moore et al.\\ 1998) had the resolution to trace the formation of small halos within galaxy clusters, with masses $\\approx 10^9$M$_\\odot$. We do not yet know what the lower mass limit is for the formation of halos in the cluster environment: determining the lower limit of galaxy mass in clusters will provide an important constraint on these models. Most of the smallest cluster galaxies are low surface brightness dwarfs for which mass estimates are very difficult, though comparison with field low surface brightness dwarfs would suggest that they may be dark matter dominated (e.g.\\ Carignan \\& Freeman 1988).  In this paper we describe a population of small objects we have found in the Fornax Cluster (see also Minniti et al.\\ 1998 and Hilker et al.\\ 1999) which have high surface brightness. The origin and nature of these objects is not yet clear, but if they are a product of the galaxy formation process in clusters, their high surface brightness will make it possible to probe the low-mass limit discussed above. They may represent extreme examples of compact low luminosity (``M32-type'') dwarf ellipticals. Alternatively, these objects may be super-massive star clusters---there is a very large population of globular clusters associated with the central galaxy of the Fornax Cluster, NGC 1399 (Grillmair et al.\\ 1994). These objects are generally similar to Galactic globular clusters with similar colours and luminosities (Forbes et al.\\ 1998). There is evidence that they are not all bound to the NGC 1399 system.  Kissler-Patig et al.\\ (1999) show that the kinematics of 74 of the globular clusters indicate that they are associated with the cluster gravitational potential rather than that of NGC 1399. They infer that the most likely origin of these globular clusters is that they have been tidally stripped from neighbouring galaxies. This has also been suggested by West et al.\\ (1995), although the effect would be diluted by the large number of halo stars that would presumably be stripped at the same time.  Bassino, Muzzio \\& Rabolli (1994) suggest that the NGC 1399 globular clusters are remnants of the nuclei of dwarf nucleated galaxies that have survived the disrupture of being captured by the central cluster galaxy. A related suggestion is a second model proposed by West et al.\\ (1995) that intra-cluster globular clusters could have formed in situ in the cluster environment.  Bassino et al.\\ (1994) conclude their discussion by noting that remnant nuclei an order of magnitude larger (and more luminous) than standard globular clusters would also be formed in significant numbers, but that existing globular cluster searches would not have included them.  In Section~\\ref{sec-obs} of this paper we describe how the observations of our {\\em Fornax Spectroscopic Survey} have sampled this part of the cluster population by measuring optical spectra of all objects brighter than $B_J=19.7$ in the centre of the Fornax Cluster. In Section~\\ref{sec-prop} we describe the properties of a new population of compact objects found in the cluster that appear to be intermediate in size between globular clusters and the smallest compact dwarf galaxies. We discuss the nature of these objects in Section~\\ref{sec-discuss} and show that higher resolution observations will enable us to determine if they are more like globular clusters or dwarf galaxies. ",
        "conclusions": "\\label{sec-discuss}  We cannot say much more about the nature of these objects on the basis of our existing data. In ground-based imaging, they are intermediate between large GCs and small compact dwarf galaxies, so it becomes almost a matter of semantics to describe them as one or the other. The most promising way to distinguish between these possibilities is to measure their mass-to-light (M/L) ratios. If they are large, but otherwise normal, GCs, they will be composed entirely of stars giving very low M/L. If they are the stripped nuclei of dwarf galaxies we might expect them to be associated with some kind of dark halo, but we would not detect the dark halos at the small radii of these nuclei, so we would also measure small M/L values. Alternatively, these objects may represent a new, extreme class of compact dwarf elliptical (``M32-type'') galaxy. These would presumably have formed by gravitational collapse within dark-matter halos, so would have high mass-to-light-ratios, like dwarf galaxies in the Local Group (Mateo 1998).  One argument against this interpretation is the apparent lack of M32-like galaxies at brighter luminosities (Drinkwater \\& Gregg 1998).  If the compact objects are dwarf galaxies, they will represent the faintest M32-like galaxies ever found. They may also fill in the gap between globular clusters and the fainter compact galaxies in the surface brightness vs.\\ magnitude distribution given by Ferguson \\& Binggeli (1994).  A further possibility is that these are small scale length ($\\sim 100$~pc) dwarf spheroidal galaxies of only moderately low surface brightness. While Local Group dSphs of equivalent luminosities generally have substantially larger scale sizes (and consequently lower surface brightnesses) (Mateo, 1998), Leo I for example has $M_B = -11.0$, and a scale length of only 110~pc (Caldwell et al., 1992), but as we discuss above this would be resolved in our existing imaging.  Our existing data will only allow us to estimate a conservative upper limit to the mass of these objects. If we say that the core radii of the objects are less than 75\\pc\\ and the velocity dispersions are less than 400\\kms\\ (the resolution of our 2dF spectra) we find that the virial mass must be less than $10^{10}\\mo$. For a typical luminosity of $M_B=-12$ this implies that $M/L < 2\\times 10^{3}$. This is not a very interesting limit, so we plan to reobserve these objects at higher spectral resolution from the ground and higher spatial resolution with the {\\em Hubble Space Telescope} (HST) in order to be sensitive to $M/L\\approx 100$. This will allow us to distinguish globular clusters from dwarf galaxies.  In order to demonstrate what we could measure with high-resolution images, we present two extreme possibilities in Figure~\\ref{fig-profile}: a very compact Galactic globular cluster and a dwarf galaxy with an $r^{1/4}$ profile ($r_e=0.2$ arcsec), both normalised to magnitudes of $B=19$ ($V=18.4$) and the Fornax cluster distance.  We also plot the PSF of the {\\em Space Telescope Imaging Spectrograph} (STIS) in the Figure for reference. The globular cluster profile is that of NGC 2808 (Illingworth \\& Illingworth 1976) with the radius scaled to the distance of the Fornax Cluster and the surface brightness then scaled to give the desired apparent magnitude. The globular cluster profile is very compact and will only just be resolved with HST, but it will clearly be differentiated from the dwarf galaxy profile.  \\begin{figure} \\hfil \\psfig{file=hst_glob_f2.eps,angle=0,width=9cm} \\caption{Predicted radial surface brightness profiles of the compact objects in two extreme cases: (A) a Galactic globular cluster (Illingworth \\& Illingworth 1976) scaled 3 mag brighter in surface brightness, and (B) a compact dwarf galaxy with an $r^{1/4}$ profile. Both are scaled to have total magnitudes of B=19 mag; they are not corrected for instrumental PSF which is also shown (C). \\label{fig-profile}} \\end{figure}  In addition to measuring the size of these objects for the mass measurement, the radial surface brightness profiles may also give direct evidence for their origin and relationship to other kinds of stellar systems. For example, if they are the stripped nuclei of galaxies, the remnants of the outer envelope might show up in the HST images as an inflection in the surface brightness profile at large radius."
    },
    "0002/astro-ph0002529_arXiv.txt": {
        "abstract": "A two-layer spherical $\\alpha^2$-dynamo model  consisting of an inner electrically conducting core (magnetic diffusivity $\\lambda_i$ and radius $r_i$) with $\\alpha = 0$ surrounded by an electrically conducting spherical shell (magnetic diffusivity $\\lambda_o$ and radius $r_o$) with a constant $\\alpha$ is shown to exhibit oscillatory behavior for values of $\\beta = \\lambda_i/\\lambda_o$ and $r_i/r_o$  relevant to the solar dynamo. Time-dependent dynamo solutions require $r_i/r_o \\geq 0.55$  and  $\\beta \\leq O(1)$. For the Sun,  $r_i/r_o$ is about 0.8  and $\\beta\\approx 10^{-3}$. The time scale  of the oscillations matches the 22   year period of the sunspot cycle for  $\\lambda_0 = O(10^2\\mbox{\\thinspace km}^2\\mbox{\\thinspace s}^{-1}$). It is unnecessary to hypothesize an $\\alpha\\omega$-dynamo to obtain  oscillatory dynamo solutions; an $\\alpha^2$-dynamo suffices provided the spherical shell region of dynamo action lies above a large, less magnetically diffusive core, as is the case for the solar dynamo. ",
        "introduction": "$\\alpha^2$- and $\\alpha\\omega$-dynamo models have been studied for decades \\citep{Bra64,Ste66,Rob72,Mof78,Gub87,Bar87}. The $\\alpha^2$-dynamo is usually stationary \\citep[e.g.,][]{Rob72,Gub87,Hol96}, although oscillatory $\\alpha^2$-dynamos have been found to occur in the special circumstance wherein $\\alpha$ changes rapidly in boundary layers \\citep{Rad87,Bar87}. In such cases, the period of the $\\alpha^2$-dynamo depends strongly on the location of the $\\alpha$-boundary layer and is typically an order of magnitude or more smaller than the magnetic diffusion time across the dynamo generation region. In general, oscillatory dynamo behavior has been produced by combination of the $\\alpha$- and $\\omega$-effects.  Kinematic models of the solar dynamo, which is inherently oscillatory, have been of the $\\alpha\\omega$-type \\citep[e.g.,][]{Roa97}.  In this paper, we report that spherical oscillatory $\\alpha^2$-dynamos can be simply induced by the magnetic coupling between an electrically conducting outer fluid shell and a conducting inner spherical core even when $\\alpha$ in the outer shell is a constant.  The period of oscillation is of the same order of magnitude as the magnetic diffusion time across the outer shell and depends largely on the electrical conductivity of the inner core.  Oscillatory behavior occurs when the outer region of dynamo action surrounds a large, less magnetically diffusive core. The radiative interior of the Sun is a large region with a smaller magnetic diffusivity than the overlying convection zone wherein dynamo action occurs. The oscillatory character of the Sun's magnetic field, as expressed in the 22 year periodicity of the sunspot cycle, could then be related to the electromagnetic coupling of the region of magnetic field generation in the convection zone with the radiative core of the Sun. The importance of an inner electrically conducting core to the problem of magnetic field generation in an overlying spherical shell has been emphasized in the $\\alpha^2$-type models of the geodynamo by \\citet{Hol93,Hol95}, \\citet{Hol96}, and \\cite{Gub99}. Nevertheless, the effects of an inner core, with magnetic diffusivity different from that of the overlying fluid convecting shell in which dynamo action takes place, have not been fully elucidated.   The problem of inner core-fluid shell coupling is made difficult by complicated electromagnetic matching conditions at the interface between the regions. For this reason, and to facilitate understanding of the physical effects, we consider the simplest type of $\\alpha^2$-dynamo model consisting of a spherical shell with $\\alpha = \\mbox{constant}$ surrounding a core with $\\alpha = 0$. We derive the appropriate matching conditions for a core and shell of arbitrary magnetic diffusivity. These matching conditions do not appear to have been considered in previous studies of spherical $\\alpha^2$-dynamos and they result in oscillatory dynamo solutions. ",
        "conclusions": "The problem solved above is a classically simple one of the type considered by \\citet{Ste66} and \\citet{Rob72} decades ago.  Yet the effects of an inner electrically conducting core on the $\\alpha^2$-dynamo are subtle and not heretofore appreciated. They enter through the complicated electromagnetic matching conditions at the interface between the core and the surrounding shell in which dynamo action occurs. The main effect of the core is to introduce time dependence into the dynamo solutions for cores whose radii are greater than about 0.55 of the outer radius of the shell. An additional requirement for time dependence is that the core be a reasonably good electrical conductor; in terms of the magnetic diffusivity ratio $\\beta = \\lambda_i/\\lambda_o$, $\\beta\\leq O(1)$ suffices for oscillatory dynamo behavior.   The importance of all this to the solar dynamo is that the parameters of the solar dynamo satisfy the requirements of oscillatory $\\alpha^2$-dynamo solutions.  For the solar dynamo $\\eta$ is about 0.8 and $\\beta$ is about $10^{-3}$ \\citep{Mof78}. In addition, if $\\omega$ (from Figure 1) is made dimensional using the time scale $\\left(r_o - r_i\\right)^2/\\lambda_o$ with $r_o - r_i = 1.4 \\times 10^5~\\mbox{km}$ and $\\lambda_o = O\\left(10^2\\mbox{\\thinspace km}^2\\mbox{\\thinspace s}^{-1}\\right)$ (eddy magnetic diffusivity), then the period of the oscillatory dynamo solution is comparable to the 22 year period of the sunspot cycle. Thus, $\\alpha^2$-dynamo action alone could be responsible for the observed time dependence of the large scale solar magnetic field. It is not our intent to suggest that the $\\omega$-effect is not significant in dynamo action in general, or in the solar dynamo in particular, because it represents a physically important process.  Our purpose is only to clarify some physics and demonstrate the potential importance of a hitherto overlooked effect, that of the oscillatory $\\alpha^2$-dynamo.  A detailed analysis of the cases $\\alpha = \\alpha\\/(r)$ and $\\alpha\\propto \\cos \\theta (\\theta = \\mbox{polar\\ angle})$ is in progress."
    },
    "0002/gr-qc0002044_arXiv.txt": {
        "abstract": "The recent analysis of Markovic and Shapiro on the effect of a cosmological constant on the evolution of a spherically symmetric homogeneous dust ball is extended to include the inhomogeneous and degenerate cases. The histories are shown by way of effective potential and Penrose-Carter diagrams. ",
        "introduction": "Recently, Markovic and Shapiro \\cite{MS}, motivated by some observational suggestions of a positive cosmological constant \\cite{Lambdaobs}, have reexamined the effect of this constant on the evolution of a homogeneous dust ball embedded in vacuum. This paper extends their analysis so as to include the inhomogeneous and degenerate cases. The qualitative behavior of the boundary histories are shown by way of effective potential and Penrose-Carter diagrams. The case $\\Lambda < 0$ is included as it provides for an interesting contrast. The well known case $\\Lambda = 0$ is not included. ",
        "conclusions": "The qualitative history of the geodesics of (\\ref{vacuum}), and via (\\ref{gammab}) therefore of the dust boundary $\\Sigma$, can be obtained from a sketch of $P$ (and in particular the requirement that $\\gamma^2 \\geq P$). These are shown in Figure 1. (The roots $(r_0,r_2,r_3)$ are given explicitly in \\cite{lake2}.) Note that for $\\Lambda < 0$ \\textit{all} orbits are closed, in contrast to $\\Lambda \\geq 0$. The case $\\gamma = 0$ is unique in the sense that $\\Sigma$ traverses the bifurcation of the Killing horizons (in the non-degenerate cases).  The Penrose - Carter diagrams are shown in Figure 2. The possible histories of $\\Sigma$ are shown. The dust can be matched to the left or to the right. The degenerate case $3 m = 1/\\sqrt{\\Lambda}$ requires a special coordinate construction \\cite{lake3}. Note that here the case $\\gamma = 0$ is associated with unstable equilibrium at the points of internal infinity.  \\begin{figure} \\caption{ Sketches of the effective potential $P$ for timelike geodesics. Choose $\\gamma^2 \\geq 0$ and note that $\\gamma^2 \\geq P$. $\\overline{r}\\ \\equiv (\\frac{3 m}{\\Lambda})^{1/3}$. The convention for the Lagrangian is $2L = -1$. } \\end{figure}  \\begin{figure} \\caption{Penrose - Carter diagrams for the trajectories shown in Figure 1. In all cases the choice of future can be reversed.} \\end{figure}   \\subsection*"
    },
    "0002/astro-ph0002235_arXiv.txt": {
        "abstract": "Among the 17 volumes of results from the Hipparcos space mission, two are dedicated to variable stars. These two volumes arose from the work of two groups, one at the Geneva Observatory and one at RGO (Royal Greenwich Observatory), on the 13 million photometric measurements produced by the satellite for 118~204 stars.  The analysis of photometric time series permitted us to identify several instrumental and mathematical problems: overestimation of the precision, offsets in the zero-points depending on the field of view, mispointing effects, image superpositions, trends in the magnitudes, binarity effects (spurious periods and amplitudes) and time-sampling effects. In this article we summarize some of the problems encountered by the Geneva group. ",
        "introduction": "The variable star analysis of the Hipparcos photometric data was an iterative process in interaction with the FAST and NDAC data reduction consortia. We started with data extracts from FAST, then from NDAC and finally worked on the whole merged data set, that is on the 118~204 time series.  The result of variable star analysis is a beautiful by-product of the mission and it was clear that it had to be published at the same time as the astrometric results (ESA 1997). The time available for the photometric analysis was then short in order to match the deadlines. The teams were put under strong pressure. The approach was then to produce a robust analysis restricting the analysis to statistically well-confirmed variables, leaving suspected variables and ambiguous cases for further analysis.  Because several instrumental problems were identified and solved, the variable star study definitely improved the overall quality of the Hipparcos photometry available now on the CD-ROMs. ",
        "conclusions": "Performing accuracy photometry in space is not free from problems. The same is true for the data analysis. Once the origin of the encountered problems is identified, it is possible to cope with them and determine precisely the domains of validity of the algorithms for search of periods and amplitudes. Globally the ratio quantity-quality generated by this mission for the study of variability has no equivalent up to now."
    },
    "0002/astro-ph0002221_arXiv.txt": {
        "abstract": "In 1991, a long term $UBV$-photometry campaign of four Wolf-Rayet stars was started using the 60 cm telescope of the National Astronomical Observatory Rozhen, Bulgaria. Here we report on our observational results and discuss the light variations.  The star WR\\,137 was observed during 1991 - 1998. No indications of eclipses were found, though random light variations with small amplitudes exist, which are probably due to dynamical wind instabilities.  WR\\,140 was also monitored between 1991 and 1998. In 1993, a dip in the light curve in all passbands was observed shortly after periastron passage, with amplitude of 0.03\\,mag in $V$. This is interpreted in terms of an ``eclipse'' by dust condensation in the WR-wind. The amplitude of the eclipse increases towards shorter wavelengths; thus, electron scattering alone is not sufficient to explain the observations. An additional source of opacity is required, possibly  Rayleigh scattering. After the eclipse, the light in all passbands gradually increased to reach the ``pre-eclipse'' level in 1998. The very broad shape of the light minimum suggests that a dust envelope was built up around  the WR-star at periastron passage by wind-wind interaction, and was gradually dispersed after 1993.  Our observations of WR\\,148 (WR + c?) confirm the 4.3 d period; however, they also show additional significant scatter. Another interesting finding is a long-term variation of the mean light (and, possibly, of the amplitude) on a time scale of years. There is some indication of a 4 year cycle of that long-term variation. We discuss the implications for the binary model.  Our photometry of WR\\,153  is consistent with the quadruple model of this star by showing that both orbital periods, 6.7 d (pair A) and 3.5 d (pair B), exist in the light variations. A search in the HIPPARCOS photometric data also reveals both periods, which is an independant confirmation. No other periods in the light variability of that star are found. The longer period light curve shows only one minimum, which might be due to an atmospheric eclipse; the shorter period light curve shows two minima, indicating that both stars in pair B are eclipsing each other. ",
        "introduction": "\\begin{table*} \\caption[]{Summary of data for the program stars. The last two columns contain\\\\ the emission line-contribution to the colours (Pyper 1966)} \\label{summary} \\begin{tabular}{rrcrrrr} \\hline\\noalign{\\smallskip} WR & HD & Spectral & Comparison & Check & $\\delta(B-V)$ & $\\delta(U-B)$ \\\\ &    & types    & star, HD   & star, HD & [mag] & [mag] \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 137  & 192641 & WC7+abs  &  192538 & 192987 & +0.07 & $-0.11$ \\\\ 140  & 193793 & WC7+O4-5 &  193888 & 193926 & +0.05 & $-0.12$ \\\\ 148  & 197406 & WN7+c?   &  197619 & 196939 & +0.01 & $-0.01$ \\\\ 153  & 211853 & 2$\\times$WN+O, or & 211430 & $---$ & +0.02 & $-0.02$ \\\\ &      & WN+O and O+O &        &       &       &         \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table*}   Photometric studies of Wolf-Rayet (WR) stars during the past decades (e.g. Moffat \\& Shara 1986; Lamontagne \\& Moffat 1987; van Genderen et al. 1987; Balona et al. 1989; Robert et al. 1989; Gosset et al. 1990; Antokhin et al. 1995; Marchenko et al. 1998a; Marchenko et al. 1998b) have revealed light variations of several per cent (up to 0.1 mag) on time-scales (typically) of days. WR stars are generally believed to be evolved Population I stars, descendants of Of-type stars (Maeder 1996). They exhibit strong, dense winds (mass loss rates of $10^{-5}$  to  $10^{-4}$ M$_{\\sun}$ yr$^{-1}$) which, in most cases, hide the stellar surface. The wind-flow is dependant on time. Moffat et al. (1988, 1994) and Robert (1994) discovered the existence of small, outward moving wind condensations, which they called propagating blobs. Unlike most O-type stars, the continuum light of many WR stars originates from a layer in the dense wind ($\\tau = 1$), a ``pseudo-photosphere'' (van Genderen et al. 1987), which could be inhomogeneous because of dynamical wind instabilities. The brightness variations of some WR stars proved to be periodic and are possibly due to binary or rotation effects. Core (photospheric) eclipses as well as atmospheric eclipses have been observed. The latter are characterized by only one V-shaped minimum on the light curve, which is caused by the atmospheric eclipse of an O-type star by the WR star's extended wind (Lamontagne et al. 1996). Random light variations are common in WR stars and they are often superimposed on the regular (binary) variations, increasing the ``noise'' and sometimes even totally disturbing the underlying regular light variations. Marchenko et al. (1998b) suggested that random light variations (light scatter) may be caused by short-lived, core-induced, multimode fluctuations, propagating in the wind. Other causes of variability, such as radial pulsations (Maeder 1985) non-radial pulsations (Vreux 1985, Antokhin et al. 1995, Rauw et al. 1996) and axial rotation (Matthews \\& Moffat 1994) have been proposed for WR stars. Occasional ``eclipses'' caused by dust formation in late-type WC stars have been studied by Veen et al. (1998).   WR\\,137 is a well known dust maker (Williams 1997; Marchenko et al. 1999). However little is known about long-term light variations for that star and its binary status is still uncertain. WR\\,140 is another repeating dust maker (Williams 1997). The orbit is well determined. Because of its high eccentricity ($e = 0.84$)  the strongest wind interaction occurs at periastron passage. During the last periastron passage in 1993, WR\\,140 received much attention and has been studied at different wavelengths from X-ray to radio. However, only a few photometric studies in the optical were carried out so far and the long-term behaviour of that star is not known. Both stars WR\\,137 and WR\\,140 are included in the infrared study by Williams et al. (1987a) and reported to have dust shells.  WR\\,148 is a good candidate for a WR + c (WR plus compact companion) binary. There is some controversy about the light variations concerning the period and the shape of the light curve. Marchenko et al. (1998a) were not able to detect the 4.31 d binary period in the HIPPARCOS photometry data, otherwise well known from ground-based observations (Marchenko et al. 1996). The very ``noisy'' light curve and unusually broad minimum need further investigation.  WR\\,153 is a quadruple system (Massey 1981), containing a WN + O and an O + O system, or two WN + O pairs (Panov \\& Seggewiss 1990). During the past  years, several photometric studies have been carried out. Yet the light variability of the two pairs could not always be unambiguously separated (Lamontagne et al. 1996). Our aim is to try to solve some of these controversial questions. ",
        "conclusions": "Our photometry of WR\\,137 reveals only small amplitude ($\\le$ 0.03 mag in the $V$ passband) random light variations. No periodicity could be found. These variations should be attributed to the continuum and they are probably due to dynamical wind instabilities.  WR\\,140 exhibited remarkable light variations and a shallow light dip is seen in all passbands shortly after periastron passage in 1993.  The light attenuation lasted until 1997 or even 1998, probably because of a dust envelope built around the WR star by wind-wind interaction at periastron passage. The dust envelope was gradually dispersed. From the wavelength dependence of the light attenuation, we find strong evidence for Rayleigh (Mie-like) scattering, contributing to the opacity, in addition to electron scattering.  For WR\\,148, our photometric study confirms the 4.317364 d light variation, but reveals occasional scatter, disturbing the light curve shape. On one occasion, we see a flare-like event at a phase when the companion is in front. Our photometry reveals long-term variations of the mean light and, possibly, of the amplitude of the regular variation. There is some evidence for a periodicity of the long-term light variation and a 4 year cycle cannot be ruled out.  Our photometry of WR\\,153 is consistent with the quadruple model for this star and both the 6.6884 d and the 3.4696 d periods are seen in the light. In pair A (6.6884\\,d period) we found evidence for an atmospheric eclipse, in agreement with the results of other works, while in pair B (3.4696 d period) the eclipse is probably photospheric (core eclipse)."
    },
    "0002/astro-ph0002017_arXiv.txt": {
        "abstract": "`Wakes' of X-ray emission have now been detected trailing behind a few (at least seven) elliptical galaxies in clusters.  To quantify how widespread this phenomenon is, and what its nature might be, we have obtained a deep (70~ksec) X-ray image of the poor cluster Abell~160 using the \\emph{ROSAT}~HRI.  Combining the X-ray data with optical positions of confirmed cluster members, and applying a statistic designed to search for wake-like excesses, we confirm that this phenomenon is observed in galaxies in this cluster.  The probability that the detections arise from chance is less than $3.8\\times10^{-3}$.  Further, the wakes are not randomly distributed in direction, but are preferentially oriented pointing away from the cluster centre.  This arrangement can be explained by a simple model in which wakes arise from the stripping of their host galaxies' interstellar media due to ram pressure against the intracluster medium through which they travel. ",
        "introduction": "\\label{sec:introduction}  With the advent of satellite-based X-ray astronomy, it was discovered that elliptical galaxies can contain as much interstellar gas as their spiral kin [for example, see Forman et al.\\ (1979)].  In the case of ellipticals, however, this interstellar medium (ISM) predominantly takes the form of a hot plasma at a temperature of $\\sim 10^7\\,{\\rm K}$.  The vast majority of elliptical galaxies are found in clusters, which themselves are permeated by very hot gas --- the intracluster medium (ICM) --- at a temperature of $\\sim 10^8\\,{\\rm K}$.  The existence of these two gaseous phases raises the question of how they interact with each other.  The collisional nature of such material means that one might expect ram pressure to strip the ISM from cluster members.  But since the ISM is continually replenished by mass loss from stellar winds, planetary nebulae, and supernovae, it is not evident \\emph{a priori} that galaxies will be entirely denuded of gas by this process.  Observing the X-ray emission from individual cluster galaxies is quite challenging, since they are viewed against the bright background of the surrounding ICM.  Sakelliou \\& Merrifield (1998) used a deep \\emph{ROSAT} observation to detect the X-ray emission from galaxies in the moderately-rich cluster Abell~2634.  They showed that the level of galaxy emission is consistent with the expected X-ray binary content of the galaxies, and hence that there is no evidence of surviving ISM in this rich environment.  When one looks in somewhat poorer environments, one does see evidence for surviving interstellar gas, and for the stripping process itself. The best-documented example is M86 in the Virgo Cluster which, when mapped in X-rays, reveals a tail or plume of hot gas apparently being stripped from the galaxy by ram pressure [Rangarajan et al.\\ (1995), and references therein].  A similar process appears to be happening to NGC~1404 in the Fornax Cluster, which displays a clear wake of X-ray emission pointing away from the cluster centre (Jones et al.\\ 1997).  Evidence for a `cooling wake' formed by gravitational accretion is found in the NGC~5044 group where a soft, linear X-ray feature is seen trailing from NGC~5044 itself (David et al. 1994).  Given the difficulty of detecting such faint wakes against the bright background of the ICM, it is quite possible that this phenomenon is widespread amongst galaxies in poor clusters.  This possibility is intriguing, since wakes indicate the direction of motion of galaxies on the plane of the sky, and it has been shown that this information can be combined with radial velocity data to solve for both the distribution of galaxy orbits in a cluster and the form of the gravitational potential (Merrifield 1998).  The existing isolated examples do not tell us, however, how common wake formation might be amongst cluster galaxies: although the observed wakes might represent the most blatant examples of widespread on-going ISM stripping, the galaxies in question might have merged with their current host clusters only recently, or be in some other way exceptional.   \\begin{figure*} \\begin{center} \\psfig{file=figures/fig1.eps,width=17.6cm,angle=0} \\end{center} \\caption{Greyscale \\emph{ROSAT} HRI image of Abell~160, smoothed using a Gaussian kernel with a dispersion of eight arcseconds. The positions of the 35 cluster galaxies with measured redshifts studied in this paper are marked with symbols such that the brightest 15 are marked with `$+$' and the fainter 20 with `$=$'. These subsamples are discussed later in the paper.} \\label{fig:overlay} \\end{figure*}   In order to obtain a more objective measure of the importance of ram pressure stripping in poor clusters, and the frequency with which it produces wakes behind galaxies, we need to look at a well-defined sample of cluster members within a single system.  The cluster Abell~160 provides an ideal candidate for such a study.  Its richness class of 0 makes it a typical poor system.  It lies at a redshift of $z=0.045$, and hence at a distance of $270\\,{\\rm Mpc}$,\\footnote{We adopt a value for the Hubble constant of $H_{0}=50\\,{\\rm km}\\,{\\rm s}^{-1}\\,{\\rm Mpc}^{-1}$ throughout this paper.} which is sufficiently close to allow galactic-scale structure to be resolved in its X-ray emission and corresponds to a size scale of $\\sim79$~kpc~arcmin$^{-1}$.  Furthermore, its Bautz-Morgan class of III means that it contains quite a number of comparably luminous galaxies, and one might hope to detect ISM emission most readily from such a sample.  In addition, its Rood-Sastry classification of C means that its members are concentrated towards the centre of the cluster:  galaxies lying in the cluster core, where the ICM density is high, will be most affected by ram pressure stripping.  Finally, Pinkney (1995) has obtained positions and redshifts for an almost complete, independently-defined set of galaxies in the field of Abell~160, providing an objective sample of cluster members for this study.  We therefore obtained a deep \\emph{ROSAT} HRI X-ray image of Abell~160 in order to investigate ISM stripping in this typical poor cluster.  We use a Galactic H\\textsc{i} column density towards A160 of $4.38\\times10^{20}$~cm$^{-2}$ (Stark et al. 1992) throughout this paper. In the next section, we present the X-ray observation and the redshift data employed.  Section~3 describes an objective method for detecting and quantifying wake features in the X-ray data, and Section~4 presents the results of applying this approach to the Abell~160 data.  We conclude in Section~5 with a discussion of the interpretation of the results.   \\begin{table*} \\centering \\caption{\\emph{ROSAT}~HRI observations of Abell~160.} \\begin{tabular}{cccc}      \\hline {\\em Observation Start Date} & {\\em Observation End Date} & {\\em Number of Used OBIs} & {\\em Total Time (s)} \\\\ \\hline 1996~Dec~30 & 1997~Jan~19 & 14 & 36965 \\\\ 1997~Jul~01 & 1997~Jul~30 & 8 & 16501 \\\\ 1997~Dec~30 & 1998~Jan~07 & 10 & 16943 \\\\ \\hline  \\\\ \\end{tabular} \\protect\\label{tab:log} \\end{table*} ",
        "conclusions": "\\label{sec:discussion}  As a first attempt at an objective determination of the frequency of wakes behind cluster galaxies, we have found significant excesses of X-ray emission apparently offset from their host galaxies.  Before exploring the possible astrophysical meaning of such features, we must rule out more prosaic possibilities.  If the X-ray emission were truly centred on the galaxies, we would still detect offset X-ray emission if there were significant positional errors in the optical galaxy locations.  The uncertainties on these positions, however, are much less than the radii at which we have detected the wakes, so this possibility can be excluded. Similarly, an overall mismatch between the optical and X-ray reference frames would produce offsets between X-ray and optical locations of coincident sources, but the distribution of apparent offset directions shown in Figure~\\ref{fig:directions} is not consistent with the coherent pattern that one would expect from either an offset or a rotation between the two frames.   \\begin{figure} \\centering \\psfig{file=figures/fig7.eps,width=8cm,angle=270} \\caption{Scatter plot of the projected angle $\\left|\\Theta\\right|$ of the 35 brightest wedge features as a function of distance from cluster centre. Larger point sizes reflect greater values of the $k$ statistic found for each galaxy, as given in the key to the figure.  } \\label{fig:scatter} \\end{figure}   We could explain the excess of sources where the X-ray emission lies at larger radii in the cluster than the optical position if the spatial scale of the optical data had been underestimated relative to that of the X-ray data.  But both the \\emph{ROSAT} and optical data image scales are extremely well calibrated.  Further, if such mismatch in magnification were responsible for the effect, one would expect the radial offsets to increase with distance from the field centre, and Figure~\\ref{fig:scatter} provides no evidence that the wakes become more radially oriented at large distances from the cluster centre.  A further possibility is that the distorted nature of the X-ray emission could arise from an asymmetry in the \\emph{ROSAT} HRI point-spread function (PSF).  Such asymmetries are documented [see, for example, Morse et al.\\ (1995)], but the observed shape of the HRI PSF actually becomes tangentially extended at large off-axis angles, so one would expect the wakes to be oriented at angles of $|\\Theta| \\sim 90^\\circ$.  The three wakes at very large radii and $|\\Theta| \\simeq70-120^\\circ$ in Figure~\\ref{fig:scatter} could well result from this phenomenon, but there is no evidence for any such effect at smaller radii.  We are therefore forced to return to trying to find an astrophysical explanation for the bulk of the observed wakes.  As discussed in the introduction, an individual galaxy can emit at X-ray wavelengths due to both its hot ISM component and its contingent of X-ray binaries. Such emission could extend to the radii where we have been searching for wakes, or appear to do so due to the blurring influence of the PSF, so we might expect some wake-like features to appear simply due to this component. Such asymmetric wake features could arise from Poisson noise on intrinsically symmetric emission, or it could reflect a real asymmetry in the emission.  For example, the emission from X-ray binaries could be dominated by one or two ultra-luminous sources in the outskirts of a galaxy, leading to an offset in the net X-ray emission.  Even the X-ray wake phenomenon that we are seeking to detect can be described as an asymmetric distortion in the normal ISM emission.  How, then, are we to distinguish between these possibilities?  Perhaps the best clue as to the nature of the detected asymmetric emission comes from the distribution of the angles at which it is detected, $\\Theta$.  As we have described above, there is a deficit of wakes pointing toward the cluster centre.  It is hard to see how such a systematic effect can be attributed to any of the more random processes such as Poisson noise on a symmetric component, or even the azimuthal distribution of X-ray binaries within the galaxy.  It therefore seems highly probable that we are witnessing the more systematic wake phenomenon that we seek.  If a wake indicates the direction of motion of the galaxy, then the deficit of detections at large values of $|\\Theta|$ implies that the production mechanism becomes ineffective when a galaxy is travelling out from the cluster centre. This conclusion has a simple physical explanation: if a galaxy is travelling on a reasonably eccentric orbit, by conservation of angular momentum it will spend a large fraction of its time close to the orbit's apocentre.  During this period, its velocity is slow and the ICM it encounters is tenuous, so it is able to retain its ISM.  In fact, continued mass loss from stellar winds and planetary nebulae means that the amount of gas in its ISM will increase. Ultimately, however, its orbit will carry it inward toward the core of the cluster.  At this point, the galaxy is travelling more rapidly, and encounters the higher density gas near the centre of the cluster, so ram pressure stripping becomes more efficient, and a wake of stripped ISM material will be seen behind the infalling galaxy.  By the time the galaxy passes the pericentre of its orbit, the ISM will have been stripped away to the extent that the outgoing galaxy does not contain the raw material to create a measurable wake, explaining the lack of detected wakes at large values of $|\\Theta|$.  This simple picture seems to fit the data on Abell~160 rather well; a similar scenario was invoked by McHardy (1979) to explain the locations of weak radio sources in clusters.  It is also notable that the beautiful wake feature behind NGC~1404 in the Fornax Cluster detected by Jones et al.\\ (1997) is oriented such that it points radially away from the cluster centre.  Clearly, though, more deep X-ray observations of clusters are required if we are to confirm the widespread applicability of this scenario.   \\subsection*{Acknowledgements} The authors are grateful to the referee for helpful comments and suggestions, and to Ian McHardy for several fruitful discussions. ND acknowledges receipt of a PPARC Studentship. This research has made use of data obtained from the Leicester Database and Archive Service at the Department of Physics and Astronomy, Leicester University, UK."
    },
    "0002/astro-ph0002151_arXiv.txt": {
        "abstract": "I use simulations of cosmological reionization to quantify the effect of photoionization on the gas fraction in low mass objects, in particular the characteristic mass scale below which the gas fraction is reduced compared to the universal value. I show that this characteristic scale can be up to an order of magnitude lower than the linear theory Jeans mass, and that even if one defines the Jeans mass at a higher overdensity, it does not track the evolution of this characteristic suppression mass.  Instead, the filtering mass, which corresponds directly to the scale over which baryonic perturbations are smoothed in linear perturbation theory, provides a remarkably good fit to the characteristic mass scale.  Thus, it appears that the effect of reionization on structure formation in both the linear and nonlinear regimes is described by a single characteristic scale, the filtering scale of baryonic perturbations.  In contrast to the Jeans mass, the filtering mass depends on the full thermal history of the gas instead of the instantaneous value of the sound speed, so it accounts for the finite time required for pressure to influence the gas distribution in the expanding universe. In addition to the characteristic suppression mass, I study the full shape of the probability distribution to find an object with a given gas mass among all the objects with the same total mass, and I show that the numerical results can be described by a simple fitting formula that again depends only on the filtering mass.  This simple description of the probability distribution may be useful for semi-analytical modeling of structure formation in the early universe. ",
        "introduction": "The effect of cosmological reionization on the formation and evolution of low mass objects has been under the scrutiny of theorists for a long time, ever since Ikeuchi (1986) and Rees (1986) independently pointed out that the increase in the temperature of the cosmic gas during reionization will suppress the formation of small galaxies with masses below the Jeans mass. Several attempts to quantify the effect of reionization on low mass galaxies have been made since using semi-analytical calculations (Babul \\& Rees 1992; Efstathiou 1992; Shapiro, Giroux, \\& Babul 1994), spherically symmetric modeling (Haiman, Thoul, \\& Loeb 1996; Thoul \\& Weinberg 1996), and three-dimensional cosmological hydrodynamic simulations (Quinn, Katz, \\& Efstathiou 1996; Weinberg, Hernquist, \\& Katz 1997; Navarro \\& Steinmetz 1997).  While confirming the general proposition that reionization suppresses formation of low mass galaxies, these studies do not give the full description of the impact of reionization on the gas fraction in the low mass objects. Particularly, one can expect that the effect of reionization depends on the reionization history, and thus is not universal at a given redshift.  Thus, it would be useful to attempt to quantify the effect of reionization on the formation and evolution of the low-mass objects in a more complete manner, relating the effective mass below which an object is a subject to reionization feedback to the characteristic scales present at each given moment of time.  \\def\\tableone{ \\begin{deluxetable}{cccccc} \\tablecaption{Simulation Parameters\\label{tabone}} \\tablehead{ \\colhead{Run} & \\colhead{$N$} & \\colhead{Box size} & \\colhead{Baryonic mass res.} & \\colhead{Total mass res.} & \\colhead{Spatial res.} } \\startdata A & $128^3$ & $4h^{-1}{\\rm\\,Mpc}$ & $10^{5.7}\\dim{M}_{\\sun}$ & $10^{6.6}\\dim{M}_{\\sun}$ &$1.0h^{-1}{\\rm\\,kpc}$ \\\\ B & $128^3$ & $2h^{-1}{\\rm\\,Mpc}$ & $10^{4.8}\\dim{M}_{\\sun}$ & $10^{5.7}\\dim{M}_{\\sun}$ &$0.5h^{-1}{\\rm\\,kpc}$ \\\\ \\enddata \\end{deluxetable} } \\placefig{\\tableone} This paper attempts to accomplish precisely that based on the new simulations of cosmological reionization. The simulations of a representative Cold Dark Matter cosmological model\\footnote{With the following cosmological parameters: $\\Omega_0=0.3$, $\\Omega_\\Lambda=0.7$, $h=0.7$, $\\Omega_b=0.04$, $n=1$, $\\sigma_8=0.91$, where the amplitude and the slope of the primordial spectrum are fixed by the COBE and cluster-scale normalizations.} were performed with the ``Softened Lagrangian Hydrodynamics'' (SLH-P$^3$M) code (Gnedin 1995, 1996; Gnedin \\& Bertschinger 1996) and fully described in Gnedin (2000). Table \\ref{tabone} lists two simulations used in this paper. Parameter $N$ gives the number of the dark matter particles; the quasi-Lagrangian baryonic mesh has the same size. Baryonic mass resolution is an average mass of a baryonic cell, and the total mass resolution is the mass of a dark matter particle plus the average mass of a baryonic cell. The spatial resolution is measured as the gravitational softening length (the real resolution of both the gravity solver and the gas dynamics solver is a factor of two worse). Reionization by stars (i.e.\\ with a soft UV background spectrum) is modeled with the Local Optical Depth approach, which is able to approximately follow the three-dimensional radiative transfer in the cosmological density distributions.  The two simulations from Table 1 allow me to investigate the sensitivity of my results to the missing small and large scale power: run A has a larger box whereas run B has a higher resolution. They also have different reionization histories, which allows me to test the generality of my results.  Both simulations have sufficient mass resolution and the box size to resolve the relevant characteristic mass scales during reionization and thus can be used for the purpose of this paper. Since both simulations have box sizes comparable to the nonlinear scale at the present time, they cannot be continued until $z=0$. Rather, run A is stopped at $z=4$ and run B at $z=6.5$.  The goal of this paper is to quantify the relationship between the total mass of an object $M_{\\rm t}$ and its gas mass $M_{\\rm g}$. The advantage of using the simulations listed in Table \\ref{tabone} is that they have enough resolution (both in mass and space) to actually map the full two-dimensional distribution of objects in the $M_{\\rm t}-M_{\\rm g}$ plane. But before this can be done, I need to discuss what characteristic mass scales are relevant for the evolution of the cosmic gas. This is particularly important because, as the reader is reminded in the next section, the Jeans mass, initially proposed as the characteristic scale, is essentially irrelevant in the expanding universe. ",
        "conclusions": "Based on cosmological simulation of reionization, I showed that the linear Jeans mass does not control the mass scales over which reionization suppresses the gas fraction in low mass cosmological objects, and may overestimate the characteristic mass scale by an order of magnitude. Since the Jeans mass is a function of the density, at the virial overdensity of 180 the Jeans mass is able to match the characteristic scale at later times, but the general shape of the Jeans mass vs redshift and the characteristic mass vs redshift do not match. Instead, the filtering mass, which directly corresponds to the length scale over which the baryonic perturbations are smoothed in linear theory, provides a good fit to the characteristic mass scale. This conclusions supports a very simple picture (proposed by Shapiro et al.\\ 1994) in which the effect of reionization on structure formation in the universe is controlled by a single mass scale both in the linear and the nonlinear regime. However, this work demonstrates that it is not the Jeans mass (as was previously thought) but the filtering mass that is the approriate mass scale.  The distribution of the gas fractions of all cosmological objects with the same total mass at any given moment during the evolution of the universe is approximately lognormal, is fully specified by the filtering mass at this moment, and is given by equations (\\ref{fitfor}-\\ref{sigfit}).  These equations, supplemented with the total-mass function of cosmological objects, fully describe the full probability distribution to find an object with given values of its gas and total mass. This probability distribution can be conveniently used in semi-analytical modeling of the evolution of low mass objects in the universe.  Admittedly, the simulations used in this paper consider only one cosmological model, but all cosmological models (including those that are curvature- or vacuum-dominated today) have an expansion law that approaches $a\\propto t^{2/3}$ at sufficiently high redshift, and follow similar behavior. Thus, while the applicability of equations (\\ref{fitfor}-\\ref{sigfit}) for a wide range of cosmological models is not rigorously proven, the physical reasoning suggests that this is likely to be the case.  My simulations also do not cover all the possible reionization histories, but at least runs A and B have different reionization histories and different resolutions, which gives some credibility to the conjecture that equations (\\ref{fitfor}-\\ref{sigfit}) work for different reionization histories and are also free from numerical artifacts."
    },
    "0002/astro-ph0002198_arXiv.txt": {
        "abstract": "One of the challenges of X-ray astronomy is how to both collect large numbers of photons yet attain high angular resolution. Because X-ray telescopes utilize grazing optics, to collect more photons requires a larger acceptance angle which in turn compromises the angular resolution. All X-ray telescopes thus have angular resolution far poorer than their diffraction limit. Although collecting more photons is a desirable goal, sometimes selective collecting fewer photons  may yield more information. Natural (such as lunar) occultations have long been used to study sources on small angular scales. But natural occulters are of limited utility because of their large angular velocities relative to  the telescope, and because of the serendipity of their transits. We describe here how one can make use of an X-ray Big Occulting Steerable Satellite (X-BOSS) to achieve very-high resolution of X-ray sources. An X-BOSS  could significantly improve the resolution of existing X-ray facilities such as the Chandra telescope, or X-ray Multiple Mirror (XMM) satellite, and could vastly improve the resolution of some future X-ray telescopes, particularly Constellation X where sub-milliarcsecond resolution is possible for a wide range of sources. Similar occulting satellites could also be deployed in conjunction with planned space observatories for other wavebands. ",
        "introduction": "One of the big challenges in doing X-ray astronomy is the relatively low photon fluxes from  target sources. The fact that X-ray mirrors operate only at grazing angles of incidence further exacerbates this problem. Thus, while one might naively expect superb angular resolution from a $1.2\\meter$ aperture X-ray telescope such as the one on board the Chandra satellite, the $0.5$ arcsecond reality is far from the 0.3 milliarcsecond nominal diffraction  limit, and considerably worse than what is routinely achieved in longer wavelength bands.  This situation is unlikely to change in the near future.  Indeed, current plans for future X-ray missions opt for increased acceptance angle (and thus increased photon count rate)  at the price of reduced angular resolution.  But it {\\em is} possible to achieve higher X-ray photon count rates and yet improve one's angular resolution. The necessary step is to separate the collection of photons from the means of achieving high resolution. One way to do this is well-known---occultation. When an astronomical body, such as the moon, transits the field of view of a telescope, it occults different sources within the field of view at different times.  By carefully measuring the photon count rate as a function of time during the transit, one can then reconstruct the projection of the surface brightness in the field of view onto the path of the occulter.   Natural occulters {\\em have} been used to achieve high-resolution in X-ray observations; however, they have at least two distinct disadvantages: \\begin{enumerate} \\item Although natural occultations can be predicted, they cannot be scheduled---target sources are therefore limited, and multiple occultations of the same source over the course of a few years are uncommon. \\item Natural occulters have large angular velocities relative to a telescope.  The shorter the transit time, the fewer photons one collects, and so the lower the resolution.  This is especially important for X-ray astronomy, where photon count rates are relatively low. \\end{enumerate} There is however an alternative to natural occulters which can overcome both of these disadvantages---a steerable occulting satellite.  Deployment of large steerable occulting satellites has been discussed for optical and near infra-red wavebands \\citep{Adams,Schneider,CS98,CS00}, mostly for the purpose of finding planet around nearby stars, but also for high-resolution astronomical observations.  However, such satellites are naturally well-suited for observations in the X-ray and far-UV\\@.  In the longer wavelength bands, minimization of diffractive losses pushes one to make the satellite as large as feasible, and deploy it as far as possible from the telescope.  In the X-ray waveband one is far into the geometric optic limit and diffraction of the X-ray photons around the satellite can essentially be neglected; thus the optimal size and placement of the satellite are determined by one's ability to accurately position the satellite with respect to the telescope-star line-of-sight and to minimize the satellite's velocity perpendicular to that line-of-sight.  The resolution delivered by the combination of the X-ray telescope and the X-BOSS is determined by the collecting area of the telescope (and thus the photon count rate for a source) and by the accuracy with which one can match the X-BOSS and telescope velocity.  It is independent of the intrinsic resolution of the telescope.  For an X-ray telescope either in an eccentric high-Earth orbit (Chandra and XMM) or at the L2 point of the Earth-Sun system (Constellation-X) we discuss in section \\ref{sec:location} where to position X-BOSS relative to the satellite.  In section \\ref{sec:blocking} of this letter we discuss the X-ray blocking efficiency of a thick film and what it implies for the required thickness of the occulter.  We also estimate in this section the required dimensions of an X-BOSS, which are determined mostly by limitations on telemetry. We discuss the steering of the X-BOSS in section \\ref{sec:steer}. In section \\ref{sec:resolve} we find the angular resolution that one obtains as a function of X-BOSS-telescope relative angular velocity, and of photon count rate. In section~\\ref{sec:skycoverage} we discuss the the sky coverage that one could obtain in each location. Application of these techniques to specific sources is discussed briefly in section~\\ref{sec:results}. Finally, section~\\ref{sec:conclude} contains the conclusions. ",
        "conclusions": "\\label{sec:conclude}  We have found that an X-BOSS used in conjunction with an X-ray telescope can lead to tremendous improvements in angular resolution. The trend of increasing the effective area of future X-ray telescopes at the expense of angular resolution (Table~\\ref{tab:xray-telescopes}) meshes perfectly with the benefits gained by including an X-BOSS in the mission.  Indeed, an X-ray telescope to be used with an X-BOSS is treated as a light bucket with all the resolving power coming from the X-BOSS occulting the source.  Thus an X-BOSS is an excellent addition to an X-ray telescope mission, particularly one at L2, such as Constellation X where sub-milliarcsecond resolution can be attained for a wide range of sources.  For the Chandra X-ray telescope we found that moderate improvements in angular resolution over an appreciable fraction of the sky can be achieved through the use of an X-BOSS\\@.  Similarly an X-BOSS employed in conjunction with XMM would provide tremendous improvements in the angular resolution that XMM could achieve allowing XMM to have angular resolution comparable to Chandra. An X-BOSS launched for use with Chandra or XMM would also provide an important test bed for the technology to be used with future missions."
    },
    "0002/astro-ph0002401_arXiv.txt": {
        "abstract": "We study the problem of gravitational lensing by an isothermal elliptical density galaxy in the presence of a tidal perturbation. When the perturbation is fairly strong and oriented near the galaxy's minor axis, the lens can produce image configurations with six or even eight highly magnified images lying approximately on a circle. We classify the caustic structures in the model and identify the range of models that can produce such lenses. Sextuple and octuple lenses are likely to be rare because they require special lens configurations, but a full calculation of the likelihood will have to include both the existence of lenses with multiple lens galaxies and the strong magnification bias that affects sextuple and octuple lenses. At optical wavelengths these lenses would probably appear as partial or complete Einstein rings, but at radio wavelengths the individual images could probably be resolved. ",
        "introduction": "The first gravitational lens to be discovered, Q~0957+561, has a simple double image configuration (Walsh, Carswell \\& Weymann 1979). It was quickly followed by the first four-image lens, PG~1115+080 (Weymann et al.\\ 1980). Together with Einstein ring lenses produced by extended sources (e.g.\\ MG~1131+0456, Hewitt et al.\\ 1988), double and quadruple lenses nearly exhaust the list of configurations in the more than 50 known strong gravitational lenses.\\footnote{For a summary, see http://cfa-www.harvard.edu/glensdata.} The only exceptions are 2016+112, which has three images and appears to require the rare and complicated situation of two lens galaxies at different redshifts (Lawrence et al.\\ 1984; Nair \\& Garrett 1997), and B~1933+503, which has ten images that can be nicely explained as a combination of three distinct sources, two of which are quadruply-imaged and one of which is doubly-imaged (Sykes et al.\\ 1998).  Theoretical studies have shown, though, that lenses with more than four images of a single source can exist. Schneider, Ehlers \\& Falco (1992) gave a mathematical analysis of caustic structures that yield additional images. Kochanek \\& Apostolakis (1988) surveyed models with two spherical galaxies at different redshifts and found that they can produce up to seven images. Witt \\& Mao (2000) found examples of models with an elliptical density galaxy and an external shear that can produce up to eight images, but they did not do a full survey of the models. Lenses with more than four images would be not only interesting to observe, but also very useful in determining the lensing mass distribution using the numerous constraints.  In this paper we present a systematic classification of the caustics and image configurations for lenses consisting of an isothermal elliptical density galaxy and a tidal perturbation. We show examples of lensing cross sections and magnification distributions for different image numbers. In a forthcoming paper we will study the observability of lenses with double, quadruple, sextuple, and octuple image configurations. The outline of this paper is as follows. In \\S 2 we review basic lens theory and the singular isothermal ellipsoid (SIE) and external shear lens models. In \\S 3 we present an analytic classification of the caustics in a lens model with an SIE galaxy and a tidal perturbation approximated as an external shear. In \\S 4 we study the stability of the caustics by adding a small core radius to the galaxy and letting the perturbation be produced by a neighboring galaxy or group. Finally, in \\S 5 we discuss some of the observational consequences and applications of our results. ",
        "conclusions": "We have studied the lensing properties of an isothermal elliptical density galaxy in the presence of a tidal perturbation to classify the caustic structures and identify different image configurations. In most cases the models have a simple caustic structure corresponding to standard 2-image and 4-image lens configurations. However, when the tidal shear has a magnitude and direction appropriate to partially cancel the effects of the galaxy's ellipticity, the caustics develop complicated swallowtail features that correspond to 6- and 8-image configurations that have not yet been observed. We gave a complete analytic treatment of the case of a singular galaxy with a perturbation modeled as an external shear, but we showed that the caustic structures are stable when one adds a small core radius or uses a more realistic treatment of the perturbation. In fact, the swallowtail caustic structures are generally bigger with the more realistic perturbation than with the shear approximation, which may enhance the likelihood of observing lenses with more than four images.  Our analysis has several observational consequences. First, sextuple and octuple lenses are likely to be rare because they require special lens configurations. In fact, they will probably be found only when there is a second galaxy close enough to the lens galaxy to provide a strong perturbation. While this situation is uncommon, it is not exceedingly rare: at least five lenses\\footnote{ B~1127+385 (Koopmans et al.\\ 1999), B~1359+154 (Rusin et al.\\ 1999), B~1608+656 (Koopmans \\& Fassnacht 1999), 2016+112 (see Nair \\& Garrett 1997 and references therein), and B~2114+022 (Augusto, Wilkinson \\& Browne 1996; Jackson et al.\\ 1998).} appear to have multiple lens galaxies within the Einstein radius. Still, to quantify the likelihood of finding a sextuple or octuple lens, it is necessary to compute the lensing cross sections for various combinations of the galaxy axis ratio $q$ and the shear amplitude $\\g$ and misalignment angle $\\tg$. Figure 10 shows cross sections for $q=0.5$, $\\g=0.22$, and three values of $\\tg$ that correspond to the cases shown in Figure 2. The configurations with more images have smaller cross sections and thus smaller probabilities for being observed. However they also have significantly higher magnifications, which means that magnification bias will be important (e.g.\\ Turner 1980; Turner, Ostriker \\& Gott 1984). Magnification bias will mitigate the effects of small cross sections to increase the likelihood of observing a sextuple or octuple lens. Clearly a realistic prediction of the lensing probabilities will require computation of cross sections and magnification bias for many combinations of $(q,\\g,\\tg)$ weighted by realistic populations of lens galaxies and perturbers. We plan to address these issues of observability in a forthcoming paper, both in the context of current surveys and of the Next Generation Space Telescope, where large numbers of lenses are expected (e.g.\\ Barkana \\& Loeb 1999).  Second, if discovered the sextuple and octuple lenses that we have described will be easy to identify because the lensed images lie approximately on a circle (see Figure 2). Any lens with more than four images that do not trace a circle must not be of the type described here. Indeed, in B~1933+503 there are ten images not on a circle, and it is thought that they must be associated with three different sources (Sykes et al.\\ 1998). In B~1359+154 there are six radio sources with four sources in a standard quadruple lens configuration, plus two sources inside the configuration whose interpretation is not clear (Myers et al.\\ 1999). Although observations and models suggest that there may be multiple lens galaxies (Rusin et al.\\ 1999), the fact that the six sources do not follow a circle suggests that lens cannot be explained by swallowtails produced by the lens galaxies.  Third, observed sextuple or octuple lenses would be very useful for constraining models of the lensing mass distribution. The ``extra'' images beyond the standard two or four would provide additional position and flux constraints. Even better, the sensitivity of the caustic structures to the lens galaxy ellipticity and the shear amplitude and misalignment angle means that the mere existence of six or more images should place strong constraints on those properties of the model. As a result, a sextuple or octuple lens could represent a wonderful ability to break the common degeneracy between the lens galaxy shape and the shear from the surrounding environment (see Keeton et al.\\ 2000a). Such a lens might even serve as the long-sought ``golden lens'' for measuring the Hubble constant $H_0$ (see Schechter 2000) provided that a time delay could be measured, although the time delay might be relatively short (a few days) because the images are all close to the critical curve.  These applications are predicated on the ability to resolve the individual images, and this ability might be limited when we include the effects of finite source size. Figure 11 shows that the source does not have to be very large before the images smear into a partial or complete Einstein ring. This happens for $R_{\\rm src}/\\bSIS \\sim 0.01$, where typically $\\bSIS \\sim 1\\arcsec$--$3\\arcsec$ for galaxy-scale lenses. Radio surveys can achieve sub-milli-arcsecond resolution with VLBI or VLBA mapping (e.g.\\ the CLASS survey, see Browne 2000 and references therein), so they should still be able to resolve individual images and thus find lenses amenable to these applications. By contrast, optical surveys (such as the Sloan Digital Sky Survey, see Gunn et al.\\ 1998 and Fischer et al.\\ 1999, or the Next Generation Space Telescope, see Barkana \\& Loeb 1999) are more likely find partial or complete Einstein rings. Still, new techniques for modeling Einstein rings show that they are very useful for constraining not only the lens model but also the intrinsic shape of the source (Keeton, Kochanek \\& McLeod 2000b).  Finally, our study may be relevant for the so-called ellipticity ``crisis,'' where lens galaxies are inferred to have larger ellipticities than the observed early-type galaxies (see Kochanek 1996 and references therein). From Figure 1, it is clear that the caustic structures are the largest when the shear is aligned with the major axis of the lensing galaxy so the ellipticity and shear act coherently; conversely, the caustic structures are the smallest when the shear is orthogonal to the major axis since the shear partially cancels the lens ellipticity. Observationally, this means that for a sample of quadruple lenses the shear may preferentially lie along the galaxy major axis; a good example is B~1422+231 (Hogg \\& Blandford 1994). The presence of an aligned shear is difficult to infer from lens models due to a degeneracy in the lens equation (see Witt 1996): naive lens models simply imply a model ellipticity equivalent to the combination of the true ellipticity and the shear. This effect may generate a bias toward larger inferred ellipticities for lens galaxies, although a full account of this bias awaits further investigations."
    },
    "0002/astro-ph0002292_arXiv.txt": {
        "abstract": "IC 4051 is a giant E galaxy on the outskirts of the Coma cluster core. Using archival \\hst\\ WFPC2 data, we derive the metallicity distribution, luminosity function, and spatial structure of its globular cluster system (GCS). The metallicity distribution derived from the $(V-I)$ colors has a mean $\\langle$Fe/H$\\rangle \\simeq -0.3$, a near-complete lack of metal-poor clusters, and only a small metallicity gradient with radius. We tentatively suggest that the GCS has two roughly equal metallicity subcomponents, one centered at [Fe/H] $\\sim 0.0$ and the second at [Fe/H] $\\sim -1.0$, although their identification is blurred by the photometric uncertainties. The luminosity distribution (GCLF) has the standard Gaussian-like form observed in all other giant E galaxies, with a peak (turnover) at $V^0 = 27.8$, consistent with a Coma distance of 100 Mpc. The radial profiles of both the GCS and the halo light show an unusually steep falloff which may indicate that the halo of this galaxy has been tidally truncated. Lastly, the specific frequency of the GCS is remarkably large:  we find $S_N = 11 \\pm 2$, at a level which rivals M87 and most others in the central cD-type category, even though IC 4051 is not a cD or brightest cluster elliptical. This galaxy exhibits a combination of GCS characteristics found nowhere else.  A formation model consistent with most of the observations would be that this galaxy was subjected to removal of a large fraction of its protogalactic gas shortly after its main phase of globular cluster formation, probably by its first passage through the Coma core.  Since then, no significant additions due to accretions or mergers have taken place, in strong contrast to the central Coma galaxy NGC 4874. ",
        "introduction": "The Coma Cluster ($d \\sim 100$ Mpc), as a rich Abell cluster, is the host environment for a huge range of E/S0 galaxies.  Its two central supergiants, NGC 4874 and 4889, are among the very most luminous galaxies known, and the cluster has many other large ellipticals scattered throughout its $\\sim 1-$Mpc core region and well beyond. The globular cluster systems (GCSs) around these galaxies are well within reach of the HST cameras, and thus give us an extraordinarily rich range of target galaxies for comparative GCS studies. Coma also has important implications for the determination of the Hubble constant $H_0$, since its large recession velocity ($\\sim 7100$ km s$^{-1}$) greatly exceeds any anticipated local peculiar motions, and it is situated at high galactic latitude ($b = +87\\fdg7$) nearly unaffected by foreground absorption,  {\\ic} is a giant E2 galaxy located $\\simeq 14'$ east in projection from the center of Coma and has no luminous neighbors. With an integrated magnitude $V_T^0 = 13.20$, it is the fifth brightest elliptical galaxy within the $1\\arcdeg$ central region of Coma. Thus, {\\ic} presents a good opportunity to study the GCS of a more-or-less ``normal'' large elliptical. \\cite{baum97} obtained the first deep WFPC2 exposures of \\ic\\ with exactly this purpose in mind, but they limited their published analysis to only the photometry from the PC1 frame on which the galaxy was centered.  However, considerably more information on the radial structure of the GCS, its specific frequency, and radial metallicity gradient are potentially available from the outer sections of the WFPC2 field. In this paper, we re-analyze the archival HST/WFPC2 frames from the \\cite{baum97} program and discuss the global properties of the \\ic\\ GCS using the entire body of data available in the raw exposures. ",
        "conclusions": "A brief summary of our findings for IC 4051 is that its GCS is (a) almost entirely metal-rich, albeit possibly with two narrowly separated subcomponents; (b) relatively compact in radial structure; and (c) a ``high specific frequency'' system despite that fact that its host galaxy is not a central giant elliptical nor one with a cD-type envelope.  Just as in Paper II for NGC 4874, we now attempt to use the integrated characteristics to reconstruct a partial history of the system. Formation scenarios for giant ellipticals tend to fall into three basic camps:  (a) ``in situ'' formation, whereby the galaxy condenses by dissipative collapse of gas clouds in its immediate vicinity, in one or more major bursts; (b) later mergers of pre-existing disk-type galaxies with both gas and stars; or (c) successive mergers or accretions of smaller gas-poor satellites.  Various combinations of these extremes are, of course, possible, and even likely.  For IC 4051, the lack of {\\it low-}metallicity clusters already places fairly strong constraints on the range of possible formation events. For example, the mechanism investigated by \\cite{cot98} -- in which an original metal-rich ``seed'' gE accretes dozens or hundreds of smaller satellites -- is unlikely, since these dwarf satellites would have brought in a population of hundreds or even thousands of low-metallicity clusters, which we do not see.  Similarly, merger-formation models in which gas-rich disk galaxies combine to build a composite elliptical (\\cite{ash92}) would predict a strong component of metal-poor clusters in the resulting MDF from the globular clusters that were present in the pre-merger galaxies.  These merger models also have severe difficulty in generating high specific frequency products, since increasing the cluster population {\\it relative to the field stars} by a large enough amount to produce high $S_N$ requires very large ($> 10^{10} M_{\\odot}$) input gas masses, more than is routinely available in disk galaxies today. The normal merger route does appear to be quite effective as a logical source for low$-S_N$ field ellipticals (see \\cite{har99} or \\cite{whi95a} for much more extensive discussion).  However, if either the merger or accretion processes are taken to an extreme form in which {\\it the merging objects are almost completely gaseous}, then they become closely similar to the {\\it in situ} route, and the conundrum of the missing low-metallicity clusters can be more easily circumvented. If the gas supply -- however it was assembled -- underwent most or all of its star formation in the high-pressure, high-density environment of the protoelliptical, then the conversion of gas to stars would have run much further to completion and built up the metallicity to the high levels that we now observe. Later gaseous mergers are, of course, not ruled out:  the central corotating disk in the core of IC 4051 (Mehlert \\etal\\ 1998) \\nocite{meh98} with its very high metallicity is a likely signature of such an event, though at its $\\ltsim 1$\\% contribution to the present-day luminosity, it probably did not form more than a few dozen globular clusters along with it, and even these would have mostly disrupted by now if they resided in the central few kpc of the core.  The relatively compact structure of the galaxy may be the result of tidal trimming (``harrassment'') from the Coma potential well (e.g., \\cite{moo96}). The radial velocity of IC 4051 (4940 km s$^{-1}$) is almost two standard deviations away from the Coma centroid (6850 km s$^{-1}$; see \\cite{col96}), indicating that this galaxy oscillates back and forth through the cluster and is now passing through the dense Coma core at high speed.  These elements of an evolutionary scenario for IC 4051 are in strong contrast to NGC 4874, for which we argued (Paper II) that a large fraction of its clusters (which are almost entirely low-metallicity) could have been acquired by accretions of smaller satellites. In IC 4051, we are forced to argue that the bulk of its clusters formed {\\it in situ}.   The globular clusters in these two galaxies provide unique evidence for the view that large E galaxies can form by radically different evolutionary routes.  One of the most challenging elements of IC 4051 to interpret is certainly the high specific frequency of its GCS. In the previous literature (\\cite{har91}, 2000; \\cite{bla97}, 1999; \\cite{har98}; \\cite{mcl99}) it has become conventional to associate high $S_N$ with giant galaxies at the centers of rich clusters. These central BCG's or cD's can have had histories of star and cluster formation through inflowing gas clouds and filaments, mergers, and accretions (e.g., \\cite{dub98}) that were much more extended than for normal outlying ellipticals.  Recently, the view has been developed that such high$-S_N$ galaxies should be regarded not as ``cluster-rich'' but rather as ``star-poor'' (\\cite{bla97}, 1999; \\cite{har98}; \\cite{mcl99}).  In this scheme, we postulate that the protogalactic gas started forming globular clusters at early times at a normal efficiency rate, but was then disrupted (perhaps by supernova-driven galactic winds, or by tidal shredding during infall; cf.~the papers cited above) before its star formation could run to completion.  The leftover gas now remains around these galaxies as their hot X-ray halos.  This picture, however, assumes that the globular clusters form earlier than the bulk of the field stars in any given round of star formation -- not an implausible requirement given the bulk of the observational evidence for starburst systems (see Harris 2000) and given that globular clusters emerge from the densest, most massive protocluster clouds.  McLaughlin (1999) \\nocite{mcl99} defines a globular cluster formation efficiency, measured empirically as the mass ratio $$ \\epsilon = {M_{cl} \\over {M_{\\star} + M_{gas}} } $$ where $M_{\\star}$ and $M_{gas}$ are the masses within the galaxy in the form of visible stars and in the X-ray gas respectively.  He finds that $\\epsilon$ is essentially identical in the well studied Virgo and Fornax systems M87, NGC 4472, and NGC 1399 (despite their very different $S_N$), providing evidence for a ``universal'' globular cluster formation efficiency $\\epsilon \\simeq 0.26$\\% relative to the {\\it initial protogalactic gas supply}.  The total mass ratio $\\epsilon$ is a more important indicator of cluster formation than $S_N$, which is only a measure of the cluster numbers (or equivalently total mass) relative to the galaxy light.  In other words, $S_N$ is a measure of only the gas mass $M_{\\star}$ that got converted to stars. Additional support for the near-universality of $\\epsilon$ in several other BCG's is found by Blakeslee (1999) \\nocite{bla99}.  In this view, {\\it any high$-S_N$ galaxy should then be surrounded by a massive X-ray gaseous component} whether or not it is a centrally dominant giant. Notably, IC 4051 is indeed one of the few Coma ellipticals with an individually detected X-ray halo. Dow \\& White (1995), \\nocite{dow95} from ROSAT observations of the Coma core region, find that IC 4051 is detectable at the $2-\\sigma$ level in the soft X-ray range $0.2 - 0.4$ keV, but not in the higher $0.4 - 2.4$ keV range. If it were at the $\\sim 6.3$ times closer distance of Virgo, IC 4051 would have a total $L_X \\simeq 5 \\times 10^{41}$ erg s$^{-1}$.  This level makes it quite comparable with the Virgo giant NGC 4472, which has $L_X \\simeq 6 \\times 10^{41}$ erg s$^{-1}$ in the soft X-ray regime (\\cite{fab92}; \\cite{irw96}; \\cite{mat97}; \\cite{buo98}).  However, this amount of X-ray gas corresponds to only $\\sim 5$\\% of the stellar mass (\\cite{mcl99}) and NGC 4472, as expected, has only a ``normal'' specific frequency level $S_N \\simeq 5$.  With our adopted distance ratios for Virgo and Coma, we find that IC 4051 is about half as luminous as NGC 4472, so if it has a roughly similar amount of X-ray gas mass, this gas would only make up $\\sim 10$\\% of its stellar mass.  Along with $S_N \\simeq 11$, we find that these parameters convert to a {\\it present-day} value for the mass ratio in IC 4051 of $\\epsilon \\sim 0.005$, twice as large as McLaughlin's (1999) \\nocite{mcl99} fiducial value.  Nominally, it therefore seems that IC 4051 acts against the paradigm of a universal globular cluster formation efficiency. An obvious possibility, however, is that IC 4051 originally did possess much more gas shortly after its main era of globular cluster formation, but that most of this unused material was quickly stripped away as IC 4051 went through its first few passages of the Coma core. This gas would have joined the general reservoir of hot gas spread throughout the Coma potential well. The same mechanism which resulted in this galaxy's compact structure might then have plausibly left it with the unusual combination of high $S_N$ and modest amount of X-ray gas that we now see.  A situation which would act much more strongly to falsify McLaughlin's case for a universal $\\epsilon$ would be the opposite one: that is, a galaxy with a massive X-ray halo but a ``normal'' or subnormal $S_N \\ltsim 4$.  In  such a case it would be much harder to avoid the conclusion that the formation efficiency of globular clusters was genuinely different (and low). Does the central Coma giant NGC 4874 present us with such a case? As we found in Paper II, NGC 4874 is {\\it not} a high-$S_N$ system and is embedded within a very massive X-ray envelope.  This X-ray gas is, however, so extended that must belong to the general Coma potential well as a whole, with no detectable concentrated component that can be associated with NGC 4874 itself (\\cite{dow95}).  Thus there are ambiguities in the interpretation that are hard to circumvent.  Better candidates would be E galaxies with massive X-ray halos that are not at the centers of rich clusters.  Finally, we may compare the interesting case of IC 4051 with that of its Coma neighbor NGC 4881 (\\cite{baum95}), a giant E galaxy of similar location, size, and structure.  Curiously, NGC 4881 holds a GCS of {\\it low} specific frequency ($S_N \\ltsim 2$) which appears to be almost entirely metal-{\\it poor}, just the opposite of IC 4051.  It has no significant amounts of X-ray gas (\\cite{dow95}). We speculate that NGC 4881 may have resulted from the merger of smaller galaxies in which these metal-poor globulars had already formed.  These mergers should have been rather gas-poor to prevent the formation of newer and more metal-rich clusters. This is, however, an extremely sketchy interpretation, and there is an obvious problem with the much higher metallicity of the host galaxy light (how did the bulk of the giant E galaxy form at higher metallicity without leaving behind some metal-rich globular clusters?  See Paper II for additional discussion).  The Coma ellipticals clearly present a wide range of GCS characteristics that strongly challenge the array of current galaxy formation models."
    },
    "0002/astro-ph0002378_arXiv.txt": {
        "abstract": "Formulating the equations of motion for cosmological bodies (such as galaxies) in an integral, rather than differential, form has several advantages.  Using an integral the mathematical instability at early times is avoided and the boundary conditions of the integral correspond closely with available data. Here it is shown that such a least-action calculation for a number of bodies interacting by gravity has a finite number of solutions, possibly only one. Characteristics of the different possible solutions are explored. The results are extended to cover the motion of a continuous fluid. A method to generalize an action to use velocities, instead of positions, in boundary conditions, is given, which reduces in particular cases to those given by Giavalisco et al. (1993) \\markcite{G93} and  Schmoldt \\& Saha (1998) \\markcite{SS98}. The velocity boundary condition is shown to have no effect on the number of solutions. ",
        "introduction": "The present motions of cosmological objects, in particular galaxies, are functions of their past history.  In principle one might discover the shape of the past by calculating presently observed positions and motions backward.  However, in doing this we are faced immediately with two problems.  First, their velocities in the plane of the sky are not known, and their distances not known accurately; so perhaps half the information needed to start the calculation by Newton's equations is there.  Second, if the trajectories of the galaxies are to be traced back to very early times distances become very small and corresponding gravitational forces very large.  Small errors in present velocities or positions become heavily magnified, resulting in galaxies being formed at infinite speeds.  The problem is mathematically unstable, rather like trying to roll a marble to the top of a glass mountain, and requiring that it stop exactly on the summit\\footnote{Valtonen et al. \\markcite{V93} (1993) have found some possible solutions for the motion of the major galaxies in the Local Group and the Maffei 1/IC 342 Group by integrating equations of motion forward from an early time.  However, it is not clear that this method is generally applicable, and in any case requires a great deal of hunting about in parameter space; for their result, the Valtonen group integrated ten thousand situations.}.  To avoid these difficulties Peebles (\\markcite{P89}1989, \\markcite{P90} 1990, \\markcite{P94} 1994) formulated the problem in integral rather than differential form.  This traded the relative simplicity and definiteness of differential equations for the stability of the integral. The most important consideration in moving from the differential to integral form of the problem (apart from the mechanics of implementation) is the fact that, with the same boundary conditions, an integral calculation may produce several (or many) solutions.  An obvious question to answer is just how many there are.  This is something more than a purely mathematical concern.  Of course, if the numerical calculation of solutions can be guided in some way there is the potential for a large savings in computer time, and if the number of solutions is limited the search may be stopped when all are found.  Conversely, if the number of solutions is very large or infinite, the usefulness of the calculation is thrown into doubt (unless some method of selecting more probable solutions is found).  But the question is more fundamental than that, for the variational formulation of the cosmological problem corresponds closely to the limits of our knowledge. When the present radial velocities and positions on the sky of a number of bodies are specified and the Big Bang postulated, we find the end conditions are fixed; the action is determined by relevant physics.  The mathematical question is thus transformed into a cosmological one.  The subject of this study is the mathematical theory of variational calculations as applied to the cosmological problem.  That problem is defined as the determination of the motion of a number of bodies moving under gravitational interaction, with the requirement that all bodies must be at the same point (in proper coordinates) at $t=0$. Newtonian, rather than relativistic, calculations are employed throughout\\footnote{See Peebles \\markcite{P80} (1980) and Bondi \\markcite{B60} (1960) for the validity of this approach.}.  The cosmological problem may be interpreted as a rough approximation of the motion of galaxies, each galaxy simulated by a point mass interacting only through gravity.  This is the way most least-action calculations have proceeded, and is not a bad approximation considering the uncertainties in such data as distances and masses.  It would be more accurate, however, to consider the objects to represent the dark matter halos of galaxies (which as far as is known interact {\\em only} through gravity).  The point-mass approximation provides a reasonable simulation of gravitational effects, since multipole moments decay rapidly with distance (Dunn \\& Laflamme 1995 found them to be quite unimportant), and in any case the conclusions of this study are not affected by the detailed form of the gravitational potential used.  Of course, identifying whole galaxies with single bodies ignores internal structure (which may be significant in some cases) and the effects of mergers (which certainly are significant);  Dunn \\& Laflamme (1995) found some additional problems.  To address these matters one must turn to a continuous fluid formulation of the problem.  Section 6 generalizes the discrete-body results to this more complicated situation. ",
        "conclusions": "The important results of this study are as follows:  {\\em If the action for the cosmological variational problem can be written in proper coordinates and an integral of energy exists, there is one minimum solution.} Assuming Hamilton's Principle is used, there may be additional, stationary solutions, one for each number of kinetic foci, if multiple-pass trajectories exist.  There is a finite number in total, limited by possible values of energy. {\\em Solutions containing at least one approximately two-body orbit which passes through more than $180\\arcdeg$ in longitude are not minima.}  {\\em Kinetic foci are reached only after the momentum is normal to the force for some body in the system.}  In so far as a continuous mass distribution may be approximated by an arbitrarily large number of individual masses, {\\em a continuum least-action calculation also has a single minimum solution, but generally a very large number of stationary solutions.} These can be limited by setting a lower limit to the resolution of the calculation.  The specific size of this resolution may be difficult to determine.  {\\em A radial velocity, rather than a distance, can be used as an end point in a numerical variational calculation.}  Forms of the modified action required have been discovered by Giavalisco et al. (1993) and used by Schmoldt \\& Saha (1998). {\\em  Using such an endpoint has no effect on the number or character of solutions.}  {\\em Orbit-crossing is not necessarily related to the number of solutions of a continuum calculation.}"
    },
    "0002/astro-ph0002187_arXiv.txt": {
        "abstract": "Eclipse lightcurves of the dwarf nova IP Peg during the November 1996 outburst are analysed with eclipse mapping techniques to constrain the location and investigate the spatial structure of the spiral shocks observed in the Doppler tomograms (Harlaftis et~al. 1999). Eclipse maps in the blue continuum and in the C\\,III+N\\,III $\\lambda 4650$ emission line show two asymmetric arcs of $\\sim 90$ degrees in azimuth and extending from intermediate to the outer disc regions ($R\\simeq 0.2 - 0.6\\; R_{L1}$, where $R_{L1}$ is the distance from disc centre to the inner Lagrangian point) which are interpreted as being the spiral shocks seen in the Doppler tomograms. The He\\,II $\\lambda 4686$ eclipse map also shows two asymmetric arcs diluted by a central brightness source. The central source probably corresponds to the low-velocity component seen in the Doppler tomogram and is understood in terms of gas outflow in a wind emanating from the inner parts of the disc. We estimate that the spirals contribute about 16 and 30 per cent of the total line flux, respectively, for the He\\,II and C\\,III+N\\,III lines. Comparison between the Doppler and eclipse maps reveal that the Keplerian velocities derived from the radial position of the shocks are systematically larger than those inferred from the Doppler tomography indicating that the gas in the spiral shocks has sub-Keplerian velocities. We undertake simulations with the aim to investigate the effect of artifacts on the image reconstruction of the spiral structures. ",
        "introduction": "Accretion discs are widespread in astrophysical environments, from sheltering the birth of stars to providing the energetics for the most violent phenomena such as relativistic jets. Despite its general importance and although considerable effort both in observation and theory has been invested over the past decade, the structure and underlying physics of accretion discs remains poorly understood. One of the major unsolved problems concerns the nature of the anomalous viscosity mechanism responsible for the inward spiraling of the disc material (Frank, King \\& Raine 1992).  Best prospects for progress in understanding accretion discs physics are possibly found in mass-exchanging binaries such as Cataclysmic Variables (CVs). In these close binaries mass is fed to a non-magnetic ($B\\simlt 10^{5}$ G) white dwarf via an accretion disc by a Roche lobe filling companion star (the secondary). The sub-class of {\\em dwarf novae} comprises low-mass transfer CVs which show recurrent outbursts of 2--5 magnitudes on timescales of months either due to an instability in the mass transfer from the secondary or due to a thermal instability in the accretion disc which switches the disc from a low to a high-viscosity regime (Warner 1995 and references therein).  Spiral shocks have been advocated by various researchers as a possible mechanism for transport of angular momentum in accretion discs (Savonije, Papaloizou \\& Lin 1994) and may be the key, together with magnetic viscosity (Hawley, Balbus \\& Winters, 1999), in understanding the viscosity mechanism. The recent discovery of spiral shocks in the accretion disc of the dwarf novae IP~Pegasi in outburst -- from Doppler tomography of emission lines (Steeghs, Harlaftis \\& Horne 1997, 1998; Harlaftis et~al. 1999) -- confirmed the results of hydrodynamical simulations (Armitage \\& Murray 1998, Stehle 1999). The spiral shocks are produced in the outer regions of the disc by the tides raised by the secondary star. During the outburst the disc expands and its outer parts feel more effectively the gravitational attraction of the secondary star leading to the formation of spiral arms.  Here we report on the eclipse mapping analysis of the data obtained by Harlaftis et al. (1999; see there for observations and data reduction). Our goal is to confirm the existence, constrain the location and to investigate the spatial structure of the spiral shocks observed in the Doppler tomograms. Section\\,\\ref{dados} presents the data and gives details of the analysis procedures. In section\\,\\ref{sim} we present a set of simulations with the eclipse mapping method aimed to clarify the interpretation of the results of section\\,\\ref{results} in terms of real spiral shocks. A summary of our findings is given in section\\,\\ref{fim}. ",
        "conclusions": "\\label{fim}  We analyzed eclipse lightcurves of the dwarf novae IP Peg during the November 1996 outburst in order to confirm the existence, constrain the location and investigate the spatial structure of the spiral shocks observed in the Doppler tomograms. Our mais results can be summarized as follows:  \\begin{itemize} \\item Eclipse maps in the blue continuum and in the C\\,III+N\\,III emission line reveal two asymmetric arcs at different azimuth and radius from disc centre which are consistent with the spiral shocks seen in the Doppler tomograms. The arcs show an azimuthal extent of $\\sim 90\\degr$ and extend from the intermediate to the outer disc regions ($R\\simeq 0.2 - 0.6\\; R_{L1}$). The outer radius of the spirals is of the same order of the disc outburst radius ($R_{d}\\simeq 0.34\\;a\\simeq 0.6\\; R_{L1}$).  \\item The He\\,II eclipse map is composed of a central brightness source plus asymmetric arcs at different distances from disc centre. The symmetric component probably corresponds to the low-velocity component seen in He\\,II Doppler tomograms and is understood in terms of gas outflow in a wind emanating from the inner parts of the disc.  \\item The spirals contribute about 16 and 30 per cent of the total line flux, respectively, for the He\\,II and C\\,III+N\\,III lines, and 13 per cent in the continuum.  \\item The Keplerian velocities derived from the radial position of the shocks are systematically larger than those inferred from the Doppler tomography, indicating that the gas in the spiral shocks has sub-Keplerian velocities.  \\end{itemize}"
    },
    "0002/astro-ph0002464_arXiv.txt": {
        "abstract": "ROSAT High Resolution Imager (HRI) data from eight observations have been co-added to obtain an effective exposure of 230 ksec.  We have identified a number of features and regions with excess X-ray brightness over that from a circularly symmetric model of the well known hot gas component.  A prominent `spur' extends 4$^{\\prime}$ from the vicinity of knot A towards the south-west.  The brightness to the south and east of this spur is significantly higher than that to the north and west.  Excess brightness is also found to the East of the nucleus, with a local maximum centered on the eastern radio lobe 3$^{\\prime}$ from the core.  There are two well known relationships between radio and x-ray emission for radio galaxies in clusters: coincidence of emissions because the X-rays come from inverse Compton or synchrotron processes; and anti-coincidence caused by exclusion of hot gas from radio entities.  We present a radio/X-ray comparison to determine if either of these relationships can be isolated in M87.  The greatest obstacle we face is the unknown projection which affects both bands. ",
        "introduction": "Eight ROSAT/HRI observations of M87 were made between 1992Jun and 1998Jan to study the X-ray structure and variability of the core and jet (Harris, Biretta, and Junor, 1997 and 1998a).  To study the large scale X-ray features of M87, we have used these data to make an image with effective exposure of 230 ksec.  A preliminary analysis based on the data then available was presented at the Ringberg Workshop on M87 held in 1997Sep (Harris, Biretta, and Junor 1998b).  The radio map of M87 (see Owen, this volume) appears to indicate an exceedingly complex structure, for which it is difficult to make meaningful measurements of isolated features.  Generally, we need to define volumes and measure their emission properties.  Once you leave the inner (brightest) lobes and jet, this is not easily done.  Even for definable filaments and other features, the surface brightness is often a sum of emissivities from various entities along the line of sight.  The X-ray map suffers from the same problem although generally there is less fine structure than in the radio (at similar resolutions).  An additional complexity for the X-ray analysis is that we cannot be certain that the very large scale X-ray distribution from the hot gas of the Virgo cluster is circularly symmetric. ",
        "conclusions": "\\begin{itemize}  \\item We suspect that projection effects confuse the interpretation of both radio and X-ray features.  Harris et al. (1998b) suggested that the SW spur might be caused by a bow shock between the ICM and the ISM, and the resulting change in temperature and density might explain the presence of knot A in the jet.  Since the current evidence supports the notion that the spur is a thermal feature, a convincing explanation for the spur will probably await the spectral/spatial capabilities of Chandra and XMM.  \\item Many X-ray and radio features are located in the same general region, but it appears that they do not actually occupy the same volumes.  \\item There is some evidence for edge effects and cavities.  \\end{itemize}"
    },
    "0002/astro-ph0002522_arXiv.txt": {
        "abstract": "{%   The possibility of determining cosmological parameters on the basis of a wide set of observational data including the Abell-ACO cluster power spectrum and mass function, peculiar velocities of galaxies, the distribution of Ly-$\\alpha$ clouds and CMB temperature fluctuations is analyzed.  Using a $\\chi^2$ minimization method, assuming $\\Omega_{\\Lambda}+\\Omega_m =1 $ and no contribution from gravity waves, we found that a tilted $\\Lambda$MDM model with one sort of massive neutrinos and the parameters $n=1.12\\pm 0.10$, $\\Omega_m=0.41\\pm 0.11$ ($\\Omega_{\\Lambda}=0.59\\pm0.11$), $\\Omega_{cdm}=0.31\\pm 0.15$, $\\Omega_{\\nu}=0.059\\pm 0.028$, $\\Omega_b=0.039\\pm 0.014$ and $h=0.70\\pm 0.12$ (standard errors) matches  observational data best. $\\Omega_{\\nu}$ is higher for more species of massive neutrinos, $\\sim 0.1$ for two  and $\\sim 0.13$ for three species. $\\Omega_m$ raises by $\\sim 0.08$ and $\\sim 0.15$ respectively. The 1$\\sigma$ (68.3\\%) confidence limits on each cosmological parameter, which are obtained by marginalizing over the other parameters, are $0.82\\le n\\le1.39$, $0.19\\le\\Omega_m\\le 1$ ($0\\le\\Omega_{\\Lambda}\\le 0.81$), $0\\le\\Omega_{\\nu}\\le 0.17$, $0.021\\le \\Omega_b\\le 0.13$ and $0.38\\le h\\le 0.85$. Varying only a subset of parameters and fixing the others shows also that the observational data set used here rules out pure  CDM models with $h\\ge 0.5$, scale invariant primordial power spectrum, zero cosmological constant and spatial curvature at a very high confidence level, $>99.99\\%$. The corresponding class of MDM models are ruled out at $\\sim 95\\%$ C.L. It is notable also that this data set determines the amplitude of scalar fluctuations approximately at the same level as COBE four-year data. It indicates  that a possible tensor component in the COBE data cannot be very substantial.  } ",
        "introduction": "The last years of the past century are marked by huge efforts of the community of astronomers, physicists and astrophysicists devoted to determine the most fundamental parameters of our Universe, the cosmological parameters. The most important among them  are the mass densities of baryons $\\Omega_b$ (in units of the critical density) and of cold dark matter $\\Omega_{cdm}$, the neutrino rest masses  $m_{\\nu}$ and their total density $\\Omega_{\\nu}$, the value of cosmological term $\\Lambda$ (or $\\Omega_{\\Lambda}$), the Hubble constant $H_0$, the spatial  curvature parameter $\\Omega_k$ and the slopes $n$ and amplitudes $A$ of the primordial power spectra of scalar and tensor fluctuations.  The primordial ratio of the number of deuterium to hydrogen nuclei (D/H) created in Big Bang nucleosynthesis is the most sensitive measure of the cosmological density of baryons $\\Omega_b$. Quasar absorption systems give definite measurements of the primordial deuterium and the most accurate value of baryon density obtained recently in this way is $\\Omega_bh^2=0.019\\pm0.0024$ \\cite{bur99}.  The measurements of the neutrino rest mass is not so certain, unfortunately. Up-to-day we have only some indications for the range where it may be found. The oscillations of solar and atmospheric neutrinos registered by the SuperKamiokande experiment show that the difference of rest masses between $\\tau -$ and $\\mu$-neutrinos is $0.02<\\Delta m_{\\tau \\mu} < 0.08eV$ \\cite{fu98,pr98}. This also provides a lower limit for the neutrino mass,   $m_{\\nu}\\ge |\\Delta m|$ and does not exclude models with  cosmologically significant values $\\sim 1-20eV$. Therefore, at least two  species of neutrinos can have approximately equal masses in this range.  Some versions of elementary particle theories predict  $m_{\\nu _e}\\approx m_{\\nu _\\tau}\\approx 2.5eV$ and  $m_{\\nu _{\\mu}}\\approx m_{\\nu _s}\\sim 10^{-5}eV$, where  ${\\nu _e}$, ${\\nu _\\tau}$, ${\\nu _{\\mu}}$ and ${\\nu _s}$ denote the  electron, $\\tau -$, $\\mu -$  and sterile neutrinos accordingly (e.g. \\cite{dol95}).  The strongest upper limit for the neutrino mass comes from the observed large scale structure of our Universe: $\\sum_{i} m_{\\nu_i}/94{\\rm eV}\\le 0.3h^2$. Since observations give for the Hubble parameter an upper limit of $h=0.8$ one gets $\\sum_{i} m_{\\nu_i}\\le 18eV$. It is interesting to note that this  upper limit coincides roughly with the upper limit for the electron neutrino mass obtained from the supernova explosion SN1987A and tritium $\\beta$-decay experiments.  Important conclusions about measurements of matter density $\\Omega_m$ ($\\equiv \\Omega_b+\\Omega_{cdm}+\\Omega_{\\nu}$) come from the Supernova Cosmology Project and the High-z Supernova Search. In particular, the relation of observed brightness vs. redshift for SNeIa shows that distant supernovae are fainter than expected for a decelerating Universe, and, thus, more distant. This can be interpreted as an accelerated  expansion rate, or $\\Omega_{\\Lambda}>0$. The best-fit value is $\\Omega_{\\Lambda}={4\\over 3}\\Omega_m + {1\\over 3}\\pm 0.1$ (1$\\sigma$ error) and  $\\Omega_m=1$ models are ruled out at the 8$\\sigma$ level \\cite{per98}. For a flat Universe $\\Omega_m + \\Omega_{\\Lambda}=1$ ($\\Omega_k=0$) the best-fit values are $\\Omega_m=0.25\\pm 0.1$ and $\\Omega_{\\Lambda}=0.75\\pm 0.1$ \\cite{per98,rie98,bah99}.  An upper limit of $\\Omega_{\\Lambda}<0.7$ (95\\% C.L.) follows from gravitational lensing statistics\\cite{bar98,fal98}, just consistent with distant supernovas results.  Strong evidence against an open Universe can be derived from recent measurements of the position of the  first acoustic peak in the cosmic microwave background (CMB) power spectrum by the Boomerang experiment\\cite{mau99}. The $1\\sigma$ range for the curvature parameter derived from this experiment is $-0.25\\le \\Omega_k\\le 0.15$ \\cite{mel99} and the mean value is close to the flat Universe, $\\Omega_k\\approx 0$.  Currently there are a few completely independent and broad routes to the determination of the Hubble constant $H_0$. The direct experiments   can be divided into three groups: the gravitational lens time delay methods, the Sunyaev-Zel'dovich method for clusters and extra-galactic distance measurements. Almost all observations yield values of $H_0$ in the  range 50-80 km/sec/Mpc.  Other independent methods for the determination of cosmological parameters are based on  large scale structure (LSS) observations. Their advantage is that all parameters mentioned above can be determined together because the form and amplitude of the power spectrum of density fluctuations are rather sensitive to all of them. Their disadvantage is that they are model dependent. This approach has been carried out in several papers (e.g. \\cite{atr97,lin97,Teg99,bri99,nov99,phill} and references therein) and it is also the goal of this paper. The papers on this subject differ by the number of parameters and the set of observational data included into the analysis. In this paper a total of 23 measurements from sub-galaxy scales (Ly-$\\alpha$ clouds) over cluster scales up to horizon scale (CMB quadrupole) is used to determine eight cosmological parameters, namely the tilt of the primordial spectrum $n$, the densities of cold dark matter $\\Omega_{cdm}$, hot dark matter $\\Omega_{\\nu}$, baryons $\\Omega_b$ and cosmological constant $\\Omega_{\\Lambda}$, the number of massive neutrino species $N_{\\nu}$, the Hubble parameter $H_0$ and, in addition, the bias parameter $b_{cl}$ for rich clusters of galaxies. We restrict ourselves to the analysis of spatially flat cosmological models with $\\Omega_{\\Lambda}+\\Omega_m=1$ ($\\Omega_k=0$) and to an inflationary scenario without tensor mode. We also neglect the effect of a possible early reionization which could reduce the amplitude of the first acoustic peak in the CMB anisotropy spectrum.  In comparison to the companion paper\\cite{nov00} the influence of the uncertainties in the normalization of the scalar mode amplitude caused by experimental errors on determination of cosmological parameters is also taken into account  here. ",
        "conclusions": "We summarize, that the observational data of the LSS of the Universe considered here can be explained by a tilted $\\Lambda$MDM inflationary model without tensor mode. The best fit parameters are: $n=1.12\\pm 0.09$, $\\Omega_m=0.41\\pm 0.11$, $\\Omega_{\\nu}=0.06\\pm 0.028$, $\\Omega_b=0.039\\pm 0.014$ and $h=0.70\\pm 0.12$. All predictions of measurements are close to the experimental values given above and within the error bars of the data. The CDM density parameter is $\\Omega_{cdm} = 0.31\\pm0.12$ and $\\Omega_{\\Lambda}$  is moderate, $\\Omega_{\\Lambda}=0.59\\pm0.11$. The neutrino matter density corresponds to a neutrino mass $m_{\\nu}=94\\Omega_{\\nu}h^2\\approx2.7\\pm1.2$ eV. The value of the Hubble constant is close to the measurements  by Madore et al.\\cite{mad98}. The age of the Universe for this model equals 12.3 Gyrs which is in good  agreement with the age of the oldest objects in our galaxy \\cite{car99}. The spectral index coincides with the COBE prediction. The relation between the matter density $\\Omega_m$ and the cosmological constant $\\Omega_{\\Lambda}$ agrees well with the independent measurements of cosmic deceleration and global curvature based on the SNIa observation.  The 1$\\sigma$ (68.3\\%) confidence limits on each cosmological parameter, obtained by marginalizing over the other parameters, are $0.82\\le n\\le 1.39$, $0.19\\le\\Omega_m\\le 1$, $0\\le\\Omega_{\\Lambda}\\le0.81$, $0\\le\\Omega_{\\nu}\\le 0.17$, $0.021\\le\\Omega_b\\le 0.13$ and $0.38\\le h\\le 0.85$.  The observational data set used here rules out the CDM models with $h\\ge 0.5$, scale invariant primordial power spectrum $n=1$ and $\\Omega_{\\Lambda}=\\Omega_k=0$ at very high confidence level, $>99.99\\%$. Also pure MDM models are ruled out at $\\sim 95\\%$ C.L.  It is remarkable also that this  data  set determines the value of normalization constant for scalar fluctuations which approximately equals the value deduced from COBE four-year data. It  indicates that a possible tensor component in the COBE data cannot be very substantial.  The coincidence of the values of cosmological parameters obtained by different methods indicates that a wide set of cosmological measurements are correct and that their theoretical interpretation is consistent. However, we must also note that the accuracy of present observational data on the large scale structure of the Universe is still insufficient to determine a set of cosmological parameters with high accuracy."
    },
    "0002/astro-ph0002008_arXiv.txt": {
        "abstract": "We present ISO-SWS spectroscopy of NGC 1068 for the complete wavelength range 2.4 to 45$\\mu$m at resolving power $\\sim$1500. Selected subranges have been observed at higher sensitivity and full resolving power $\\sim$2000. We detect a total of 36 emission lines and derive upper limits for 13 additional transitions. Most of the observed transitions are fine structure and recombination lines originating in the narrow line region (NLR) and the inner part of the extended emission line region.  We compare the line profiles of optical lines and reddening-insensitive infrared lines to constrain the dynamical structure and extinction properties of the narrow line region. The most likely explanation of the considerable differences found is a combination of two effects. (1) The spatial structure of the NGC\\,1068 narrow line region is a combination of a highly ionized outflow cone and lower excitation extended emission. (2) Parts of the narrow line region, mainly in the receding part at velocities above systemic, are subject to extinction that is significantly suppressing optical emission from these clouds. Line asymmetries and net blueshifts remain, however, even for infrared fine structure lines suffering very little obscuration. This may be either due to an intrinsic asymmetry of the NLR, as perhaps also suggested by the asymmetric radio continuum emission, or due to a very high column density obscuring component which is hiding part of the narrow line region even from infrared view.  We present detections and limits for 11 rotational and ro-vibrational emission lines of molecular hydrogen (H$_2$). They arise in a dense molecular medium at temperatures of a few hundred Kelvin that is most likely closely related to the warm and dense components seen in the near-infrared H$_2$ rovibrational transitions, and in millimeter wave tracers (CO, HCN) of molecular gas. Any emission of the putative pc-scale molecular torus is likely overwhelmed by this larger scale emission.  In companion papers we use the SWS data to derive the spectral energy distribution emitted by the active nucleus of NGC\\,1068 (\\cite{alexander00}), to put limits on infrared emission from the obscured broad line region (\\cite{lutz00}), and discuss the continuum and its features in conjunction with SWS spectra of other galaxies (\\cite{sturm00}). ",
        "introduction": "NGC\\,1068 is one of the nearest and probably the most intensely studied Seyfert 2 galaxy. Observations in all wavelength bands from radio to hard X-rays have formed a uniquely detailed picture of this object. NGC 1068 has played a key role in the development of unified scenarios for Seyfert 1 and Seyfert 2 galaxies (\\cite{antonucci85}), in the study of molecular gas in the nuclear region of Seyferts (e.g. \\cite{myers87}; \\cite{tacconi94}), and in elucidating the importance of star formation activity coexistent with the AGN, both on larger (e.g. \\cite{telesco88}) and smaller (\\cite{macchetto94}; \\cite{thatte97}) scales. NGC 1068 hosts a prominent Narrow Line Region (NLR) that is approximately cospatial with a linear radio source with two lobes (\\cite{wilson83}). The narrow emission line region has been extensively characterized from subarcsecond clouds probed by HST (\\cite{evans91}; \\cite{macchetto94}), the $\\approx 5$ arcseconds of the NLR hosting most of the line flux (e.g. \\cite{walker68}; \\cite{shields75}; \\cite{cecil90}), and the ionization cone and extended emission line region (\\cite{pogge88}; \\cite{unger92}) extending to radii of at least 30\\arcsec\\/ (1\\arcsec\\/ = 72 pc at the distance of 14.4 Mpc, \\cite{tully88}). The velocity field is complex, with an ensemble of rapidly moving clouds dominating the inner arcseconds and a more quiescent rotation pattern prevailing at larger radii (e.g. \\cite{walker68}; \\cite{alloin83}; \\cite{meaburn86}; \\cite{cecil90}). While most of the excitation of the narrow line region and the extended emission line region is likely through photoionization by the central AGN (\\cite{marconi96}), high resolution observations suggest kinematic disturbance and possibly shock excitation of regions close to the radio outflow (e.g. \\cite{axon98}).  With ESA's {\\it Infrared Space Observatory} ISO, sensitive mid-infrared spectroscopy of AGNs became possible, with detections of a broad range of low- and high-excitation fine structure lines, recombination lines, and pure rotational lines from such sources (\\cite{moorwood96}; \\cite{sturm99}; \\cite{alexander99})). Model predictions of the mid-IR spectra of AGN had been obtained prior to ISO (e.g. \\cite{spinoglio92}), but observations were restricted by limited sensitivity and focussed primarily on the continuum emission and broad features rather than emission lines.  In this paper, we present the ISO-SWS spectra of NGC\\,1068 and draw conclusions on the structure of the narrow line region that can be obtained mainly from the comparison of optical and reddening-insensitive infrared lines, and discuss the nature of the mid-infrared molecular hydrogen emission. The ISO-SWS data of NGC\\,1068 are analysed further in several companion papers. Alexander et al. (2000) use photoionization modelling based on the ISO fine structure line set and other NLR lines to model the shape of the AGN's spectral energy distribution. Lutz et al. (2000) analyze limits on emission from the obscured broad line region. Finally, Sturm  et al. (2000) discuss continuum energy distribution and features of NGC\\,1068 in conjunction with ISO-SWS spectra of other galaxies.  Our paper is organised as follows. In  \\S 2 we discuss the ISO-SWS observations and data reduction. \\S 3 presents results and implications of the density of the narrow line region. \\S 4 uses infrared line profiles in comparison to optical ones to constrain the structure of the narrow line region. We discuss the mid-infrared molecular hydrogen emission in \\S 5 and summarize in \\S 6. ",
        "conclusions": "ISO-SWS spectroscopy provides the first detailed census of the mid-infrared spectrum of the prototypical Seyfert 2 galaxy NGC\\,1068. We have detected 36 emission lines on top of the strong AGN-heated continuum. Most lines originate in the NLR characterized by a density of $\\sim$2000\\,cm$^{-3}$.  We have compared the mid-infrared ISO line profiles with optical emission line profiles produced in the NLR. The line profiles are consistent with a model where the NLR is a combination of a highly ionized outflow and lower excitation extended emission, with extinction significantly affecting the optical line profiles. Remaining blueshift and asymmetry of the least obscured lines may reflect either intrinsic asymmetry of the NLR or an additional very high column density obscuring component.  We detect strong emission from warm molecular hydrogen, which most likely originates on the 100pc to kpc scale, and which is also probed by emission in near-infrared and millimeter wave tracers of molecular material. This emission masks any possible emission from a putative parsec-scale molecular torus.  Companion papers use the SWS data to model the spectral energy distribution of the active nucleus, to put limits on emission from the obscured broad line region, and discuss the continuum and its features."
    },
    "0002/astro-ph0002244_arXiv.txt": {
        "abstract": "We present new HST NICMOS observations of NGC 4945, a starburst galaxy hosting a highly obscured active nucleus that is one of the brightest extragalactic sources at 100 keV. The HST data are complemented with ground based \\FeII\\ line and mid--IR observations.  A 100pc-scale starburst ring is detected in \\PA, while \\Hmol\\ traces the walls of a super bubble opened by supernova-driven winds. The conically shaped cavity is particularly prominent in \\PA\\ equivalent width and in the \\PA/\\Hmol\\ ratio. Continuum images are heavily affected by dust extinction and the nucleus of the galaxy is located in a highly reddened region with an elongated, disk-like morphology. No manifestation of the active nucleus is found, neither a strong point source nor dilution in CO stellar features, which are expected tracers of AGN activity.  Even if no AGN traces are detected in the near-IR, with the currently available data it is still not possible to establish whether the bolometric luminosity of the object is powered by the AGN or by the starburst: we demonstrate that the two scenarios constitute equally viable alternatives. However, the absence of any signature other than in the hard X-rays implies that, in both scenarios, the AGN is non-standard: if it dominates, it must be obscured in all directions, conversely, if the starburst dominates, the AGN must lack UV photons with respect to X-rays.  An important conclusion is that powerful AGNs can be hidden even at mid-infrared wavelengths and, therefore, the nature of luminous dusty galaxies cannot be always characterized by long-wavelength data alone but must be complemented with sensitive hard X-ray observations. ",
        "introduction": "A key problem in studies of objects emitting most of their energy in the FIR/submm is to establish the relative importance of highly obscured Active Galactic Nuclei (AGN) and starburst activity. In particular, it is important to know if it is still possible to hide an AGN, contributing significantly to the bolometric emission, when optical to mid-IR spectroscopy and imaging reveal only a starburst component.  Several pieces of evidence suggest that most cosmic AGN activity is obscured. Most, and possibly all, cores of large galaxies host a supermassive black hole (\\ten{6}--\\ten{9}\\Mo; e.g. Richstone et al. \\cite{richstone}). To complete the formation process in a Hubble time, accretion must proceed at high rates, producing quasar luminosities ($L\\sim\\ten{12}\\Lo$). However the observed black hole density is an order of magnitude greater than that expected from the observed quasar light, assuming accretion efficiency of 10\\%, suggesting that most of the accretion history is obscured (e.g. Fabian \\& Iwasawa \\cite{fabian99}, and references therein). It is estimated either that 85\\% of all AGNs are obscured (type 2) or that 85\\% of the accretion history of an object is hidden from view.  In addition, the hard X-ray background ($>1\\KEV$) requires a large population of obscured AGNs at higher redshifts ($z\\sim1$) since the observed spectral energy distribution cannot be explained with the continua of Quasars, i.e. un--obscured (type 1) AGNs (Comastri et al. \\cite{comastri}, Gilli et al. \\cite{gilli99}). Despite the above evidence, detections of obscured AGNs at cosmological distances are still sparse (e.g. Akiyama et al. \\cite{akiyama}).  Ultra Luminous Infrared Galaxies (ULIRGs; see Sanders \\& Mirabel \\cite{sanders96} for a review) and the sources detected in recent far-infrared and submm surveys performed with ISO and SCUBA (e.g. Rowan-Robinson et al. \\cite{rowanrob}, Blain et al. \\cite{blain} and references therein) are candidate to host the missing population of type 2 AGNs. However, mid-IR ISO spectroscopy has recently shown that ULIRGs are mostly powered by starbursts and that no trace of AGNs is found in the majority of cases (Genzel et al. \\cite{genzel98}; Lutz et al. \\cite{lutz98}). Yet, the emission of a hidden AGN could be heavily absorbed even in the mid-IR. Indeed, the obscuration of the AGN could be related to the starburst phenomenon, as observed for Seyfert 2s (Maiolino et al. \\cite{maiolino95}). Fabian et al. (\\cite{fabian98}) proposed that the energy input from supernovae and stellar winds prevents interstellar clouds from collapsing into a thin disk, thus maintaining them in orbits that intercept the majority of the lines of sight from an active nucleus.  In this paper, we investigate the existence of completely obscured AGNs and the Starburst-AGN connection through observations of NGC 4945,  one of the closest galaxies where an AGN and starburst coexist. NGC 4945 is an edge-on ($i\\sim 80^\\circ$), nearby ($D=3.7\\MPC$) SB spiral galaxy hosting a powerful nuclear starburst (Koornneef \\cite{koorn}; Moorwood \\& Oliva \\cite{moorwood94a}). It is a member of the Centaurus group and, like the more famous Centaurus A (NGC 5128), its optical image is marked by dust extinction in the nuclear regions. The ONLY evidence for a hidden AGN comes from the hard X-rays where NGC 4945 is characterized by a Compton-thick spectrum (with an absorbing column density of $\\NH=5\\xten{24}\\CM\\2$, Iwasawa et al. \\cite{iwasawa93}) and one of the brightest 100\\KEV\\ emissions among extragalactic sources (Done et al. \\cite{done96}). Recently, BeppoSAX clearly detected variability in the 13-200\\KEV\\ band (Guainazzi et al., \\cite{guainazzi}).   Its total infrared luminosity derived from IRAS data is $\\sim 2.4\\xten{10}\\Lo$ (Rice et al. \\cite{rice88}), $\\sim 75\\%$ of which arises from a region of $\\le 12\\arcsec\\times9\\arcsec$ centered on the nucleus (Brock et al. \\cite{brock88}). Although its star formation and supernova rates are moderate, $\\sim 0.4\\,\\Mo\\YR\\1$ and $\\sim 0.05\\YR\\1$ (Moorwood \\& Oliva \\cite{moorwood94a}), the starburst activity is concentrated in the central $\\sim 100\\PC$ and has spectacular consequences on the circumnuclear region which is characterized by a conical cavity evacuated by a supernova-driven wind (Moorwood et al. \\cite{moorwood96a}).  The radio emission is characterized by a compact non-thermal core with a luminosity of $\\simeq 8\\xten{38}\\ERG\\S\\1$ (Elmouttie et al. \\cite{elmouttie}). It is one of the first H$_2$O and OH megamaser sources detected (dos Santos \\& Lepine \\cite{dossantos}; Baan \\cite{baan85}) and the H$_2$O maser was mapped by Greenhill et al. (\\cite{greenhill}) who found the emission linearly distributed along the position angle of the galactic disk and with a velocity pattern suggesting the presence of a $\\sim\\ten{6}\\Mo$ black hole. Mauersberger et al. (\\cite{mauersberger}) mapped the $J=3-2$ line of $^{12}$CO which is mostly concentrated within the nuclear $\\sim 200\\PC$.   We present new line and continuum images obtained with the {\\it Near Infrared Camera and Multi Object Spectrograph} (NICMOS) on-board the Hubble Space Telescope (HST), aimed at detecting AGN activity in the near-infrared. These observations are complemented by recent ground based near- and mid-IR observations obtained at the European Southern Observatory. Section \\ref{sec:obs} describes the observations and data reduction techniques. Results are presented in Section \\ref{sec:res} and discussed in Section \\ref{sec:discuss}. Finally, conclusions will be drawn in Sec. \\ref{sec:conclus}. Throughout the paper we assume a distance of 3.7\\MPC\\ (Mauersberger et al. \\cite{mauersberger}), whence 1\\arcsec\\ corresponds to $\\simeq18$\\PC.  \\begin{figure*} \\begin{center} \\begin{tabular}{cc} \\epsfig{figure=marconi_f1a.ps,width=0.45\\linewidth} & \\epsfig{figure=marconi_f1b.ps,width=0.45\\linewidth} \\\\ & \\\\ \\epsfig{figure=marconi_f1c.ps,width=0.45\\linewidth} & \\epsfig{figure=marconi_f1d.ps,width=0.45\\linewidth} \\\\ & \\\\ \\epsfig{figure=marconi_f1e.ps,width=0.45\\linewidth}    & \\hspace{14pt}\\vspace*{-14pt} \\epsfig{figure=marconi_f1f.ps,width=0.38\\linewidth} \\\\ \\end{tabular} \\end{center} \\caption{\\label{fig:cont} (a) F222M image (K band). North is up and East is left. The cross marks the location of the K nucleus and the circle represents the uncertainty on the position of the H$_2$O maser given by Greenhill et al. (1997). Units of the frame box are seconds of arc. The origin is at the nominal location of the H$_2$O maser. (b) F160W image (H). Notation as in panel (a). (c) F110W image (J). Notation as in panel (a). The black contours are from the H-K color image at 1.8, 2 and 2.2 levels. (d) F606W image (R band). Notation as in panel (a) except for the contours which are from the K band image. (e) H-K image. Symbols are as in (a). (f) Truecolor (Red=F222M, Green=F110W, Blue=F606W) image. } \\end{figure*}   \\begin{figure*}[!] \\begin{center} \\begin{tabular}{cc} \\epsfig{figure=marconi_f2a.ps,width=0.45\\linewidth} & \\epsfig{figure=marconi_f2b.ps,width=0.45\\linewidth} \\\\ & \\\\ \\epsfig{figure=marconi_f2c.ps,width=0.45\\linewidth} & \\epsfig{figure=marconi_f2d.ps,width=0.45\\linewidth} \\\\ & \\\\ \\epsfig{figure=marconi_f2e.ps,width=0.45\\linewidth} & \\hspace{14pt}\\vspace*{-14pt} \\epsfig{figure=marconi_f2f.ps,width=0.38\\linewidth} \\\\ \\end{tabular} \\end{center} \\caption{\\label{fig:line} (a) Pa$\\alpha$ image. Symbols as in Fig. 1. The black contours are from the \\HA+\\NII\\ image by Moorwood et al. (1996). (b) H$_2$ image. Black contours are from the blue ground-based \\FeII\\ image, (c) Equivalent width of Pa$\\alpha$. (d) \\PA/\\Hmol\\ image. Symbols as in Fig. 1. (e) CO index. Symbols as in Fig. 1. (f) Truecolor line image (Red=F222M, Green=\\Hmol, Blue=\\PA) image. } \\end{figure*} ",
        "conclusions": "Conclusions}  Our new HST NICMOS observations of NGC 4945, complemented by new ground based near and mid-IR observations, have provided detailed morphology of the nuclear region. In \\PA, we detect a 100pc-scale starburst ring while in \\Hmol\\ we trace the walls of a conical cavity blown by supernova driven winds. The continuum images are strongly affected by dust extinction but show that even at HST resolution and sensitivity, the nucleus is completely obscured by a dust lane with an elongated, disk-like morphology. We detect neither a strong point source nor dilution in CO stellar features, expected signs of AGN activity.  Whereas all the infrared properties of NGC 4945 are consistent with starburst activity, its strong and variable hard X-ray emission cannot be plausibly accounted for without the presence also of an AGN. Although the starburst must contribute to the total bolometric luminosity we have shown, using starburst models, that the actual amount is dependent on the star formation history. A major contribution from the AGN is thus not excluded. Irrespective of the assumption made, however, our most important conclusion is that the observed variable hard X-ray emission combined with the lack of evidence for reprocessed UV radiation in the infrared is incompatible with the \"standard\" AGN model. If the AGN dominates the bolometric luminosity, then its UV radiation must be totally obscured along all lines of sight. If the starburst dominates then the AGN must be highly deficient in its UV relative to X-ray emission. The former case clearly raises the possibility that a larger fraction of ULIRGs than currently thought could actually be AGN rather than starburst powered."
    },
    "0002/astro-ph0002072_arXiv.txt": {
        "abstract": "From 43\\,GHz VLBA observations of the pair of radio sources \\1928+738 and \\2007+777 we have demonstrated the feasibility of precision phase-delay differential astrometric techniques at millimeter wavelengths. For a pair of sources with 5\\degr\\, separation and high antenna elevations, we have shown that present astrometric models and millimeter arrays are advanced enough to model the differential phase-delay to within 2 picoseconds, less than one tenth of a phase-cycle at 43\\,GHz. The root-mean-square of the differential phase-delay residuals is dominated by the fluctuations of the atmospheric water vapor. We have determined the relative position of the observed sources with a precision twofold better than previous determinations at lower frequencies and, more importantly, largely free from ambiguous definitions of the reference point on the structure of the radio sources. Our result makes 43\\,GHz VLBI phase-delay differential astrometry an ideal tool to study the absolute kinematics of the highly variable structures of regions near the core of extragalactic radio sources. ",
        "introduction": "One of the trends in Very-Long-Baseline Interferometry (VLBI) is to augment the angular resolution of the observations in search for a more detailed view of the inner structure of extragalactic radio sources. This is effectively carried out by either observing at millimeter wavelengths (mm-VLBI) or, at cm-wavelengths, by combining ground telescopes with antennas in space. The correct interpretation of these high-resolution observations is of much relevance since they map the morphology of highly variable regions close to the central engine of AGNs. However, multi-epoch analyses directed to understand the dynamical behavior of these inner regions critically depend on the alignment of the images: no solid conclusions can be extracted without an accurate source component (i.e. core) identification. In particular, VLBI reveals that cm-wavelength components break up in complex structures with multiple features at mm-wavelengths. These compact mm-features show a strong variability, which may be the result of phenomena only seen so far in numerical simulations (G\\'omez et al. 1995). For a meaningful physical understanding of those compact features, a detailed knowledge of the (absolute) kinematics of the region is crucial. It is therefore highly desirable to extend precision differential phase-delay astrometry to mm-wavelengths.\\\\ In this Letter we demonstrate the feasibility of using phase-delay differential astrometry at 43\\,GHz. We have selected the pair of sources \\1928+738 and \\2007+777, $\\sim$5\\degr\\, apart, with flat spectra, high flux densities, and rich structures at 43\\,GHz. The astrometry analysis of these data show the advantages and possibilities of mm-wavelength differential astrometry. ",
        "conclusions": "\\begin{table}[t] \\caption[]{Contributions to the standard errors of the estimates of the coordinates of \\1928+738 minus those of \\2007+777 ( $\\delta\\Delta\\alpha$, $\\delta\\Delta\\delta$) from the sensitivity study. } \\begin{flushleft} \\begin{footnotesize} \\begin{tabular}{llccc} \\hline Effect  & & Standard      & $\\delta\\Delta\\alpha^b$    & $\\delta\\Delta\\delta$ \\\\ & & Deviation$^a$ &     ($\\mu$s)  &  ($\\mu$as)   \\\\ \\hline Statistical errors$^c$    &             & --               &  4   & 25 \\\\ Ref. point identification &             & --               &  0.6 &  1 \\\\ Station coordinates       &             &  2\\,cm           &  6   & 25 \\\\ Coordinates               & $\\alpha$    & 100\\,$\\mu$s      & 23   & 41 \\\\ of \\2007+777:             & $\\delta$    & 300\\,$\\mu$as     &  9   &  9 \\\\ Earth's pole:             & $x$         & 150\\,$\\mu$as     &  2   &  9 \\\\ & $y$         & 250\\,$\\mu$as     &  2   &  4 \\\\ UT1-UTC                   &             &  15\\,$\\mu$s      &  3   &  5 \\\\ Earth's nutation:         & $d\\psi$     & 170\\,$\\mu$as     &  0.3 &  1 \\\\ & $d\\epsilon$ &  80\\,$\\mu$as     &  1   &  2 \\\\ Ionosphere$^d$            &             & 1 TECU           &  1   &  2 \\\\ \\hline rss$^e$                   &             &                  & 26   & 56 \\\\ \\hline \\end{tabular} \\end{footnotesize}  \\begin{scriptsize} \\noindent $^a$  The standard deviation of the fixed geometrical parameters of our astrometric model (all entries but ionosphere) were taken from IERS Annual Report 1998 (1999). The 2\\,cm standard deviation of the site coordinates corresponds to each of the three coordinates for each antenna site.\\\\ $^b$  Notice that the values of $\\delta\\Delta\\alpha$ are in $\\mu$s. To convert $\\mu$s to $\\mu$as, the factor 15$\\cdot\\cos\\delta_{\\rm{QSO}\\,1928+738}$ ($\\sim$4.2 ) must be used.\\\\ $^c$ The statistical errors include the uncertainties of the tropospheric zenith delays at the nodes of the piecewise linear function used in the troposphere model (see Sect. 2).\\\\ $^d$ Standard deviation provided by the global ionospheric maps at each site. 1 TECU = 1$\\times$10$^{16}$ el\\,m$^{-2}$.\\\\ $^e$ Root-sum-square of the tabulated values. \\end{scriptsize} \\end{flushleft} \\end{table}   \\noindent From the astrometric analysis described in Sect. 3, we obtain the following J2000.0 coordinates of \\1928+738 minus those of \\2007+777 at 43\\,GHz:\\\\  \\noindent \\begin{tabular}{lll} $\\Delta\\alpha=$ & $-0^{h}\\,37^{m}\\,42\\rlap{.}^{s}503443$ & \\, $\\pm\\,0\\rlap{.}^{s}000026$\\\\ $\\Delta\\delta=$ & $-3^{\\circ}\\,54'\\,41\\rlap{.}''677208$   & \\,  $\\pm\\,0\\rlap{.}''000056$\\\\ \\end{tabular}  \\vspace*{0.2cm} \\noindent where the quoted uncertainties are overall standard errors (see Table 1), nearly twofold smaller than the standard errors corresponding to previous determinations at 5 and 8.4\\,GHz. From the comparison of the results of the sensitivity analysis displayed in Table 1 with similar sensitivity analyses at lower frequencies (G95; R99), we see that the improvement in precision comes from (i) the small contribution to the standard errors of the reference point identification in the map (dominated by image noise), as a consequence of the improvement of resolution of the maps and of the lack of ambiguity in selecting the components acting as reference, and (ii) the negligible contribution of the ionosphere, that scales down by a factor of 25 with respect to its contribution at 8.4\\,GHz. As occurs at cm-wavelengths, the quoted standard errors of the relative position are dominated by the uncertainties of the fixed parameters of the astrometric model (entries 3 to 10 of Table 1), and, in particular, by the uncertainties of the coordinates of the reference source, as expected for objects with a large angular separation (notice that this error is not frequency dependent). The comparison and interpretation of the relative position estimate at 43\\,GHz with previously reported estimates at lower frequencies will be postponed to a later publication where the comparison will be made in great detail.  \\noindent The postfit residuals of the differenced phase delays corresponding to all baselines included in our analysis are shown in Fig. 2. Note the scale of the plots, $\\pm$23\\,ps, corresponding to $\\pm$1 phase cycle. The average root-mean-square (rms) of the postfit residuals is 2\\,ps, less than one tenth of the phase cycle at 43\\,GHz. At this level of precision, the absence of systematic effects validates both the astrometric model, based on IERS standards, and the propagation medium procedures for mm-wavelength VLBI astrometry (at least for cycle times, source separations, weather conditions, and antenna elevations similar to those presented in this paper). To calibrate the quality of our procedure, we compared the residual of the differenced phase delay with similar residuals corresponding to observations at 8.4\\,GHz and 5\\,GHz made in the past (G95, G98). The rms of the residuals are about 15, 9, and 2\\,ps for the data sets at 5, 8.4, and 43\\,GHz, respectively, which expressed in equivalent-phase yield postfit residuals of roughly 30 degrees at each of the three frequencies. This similarity of the rms expressed in phase at all the observed frequencies demonstrates not only that the phase connection process is feasible at 43\\,GHz, but which is also of no less quality than at lower frequencies.\\\\ Likely, the most important contributors to the scatter of the phases at 43\\,GHz are the unmodeled variations of the refractivity of the neutral atmosphere. From the average rms of 2\\,ps of the phase residuals of Fig. 2, and assuming uncorrelated contributions from the antennas forming each interferometric pair, the average uncertainty for the single-site phase delay is $\\sqrt{2}$\\,ps. This uncertainty is in good agreement with the predictions of water vapor fluctuations on time scales of 100 seconds ($\\sim$2\\,ps) based on refractivity patterns described by Kolmogorov turbulences (Treuhaft \\& Lanyi 1987).\\\\ The importance of our result translates to VLBI phase-referencing mapping. This technique (see e.g. Lestrade \\et 1990) relies completely on the behavior of the phase of the reference source (usually a strong radio emitter) to predict the phase of the target source (usually a weak radio emitter). Beasley \\& Conway (1995) provide useful expressions for the maximum cycle time for phase-referencing with the VLBA.  Under good weather conditions, average antenna elevation of 40\\degr, and with the requirement that the rms phase between scans is $<$90\\degr, the maximum cycle time at 43\\,GHz is $\\sim$100s. This estimate should be shortened if atmospheric spatial variations from different lines of sight are considered. Actually, the facts are more favorable. For sources separated 5\\degr\\, on the sky, high antenna elevations, and good weather conditions, our results show that (i) the rms of the phases is below 90\\degr\\, throughout the experiment and does not seem to be substantially dependent of the cycle times used during our observation (100-300s); and (ii) the expected average uncertainty in interpolating the phases of one source to the epoch of the other is $\\sim$30\\degr\\,in the differenced phase. This value is not larger than usual phase errors in phase-reference mapping at cm-wavelengths (e.g. Lestrade \\et 1990). Therefore, with the proper cycle times and nearby calibrator sources, diffraction limited VLBI phase-reference images at 43\\,GHz should be possible.  \\begin{figure}[t] \\vspace*{9cm} \\special{psfile=guirado_f2.ps hoffset=-40 voffset=-54 hscale=50 vscale=50} \\caption[]{Postfit residuals of the difference phase delays at 43\\,GHz for all baselines. Full vertical scale is $\\pm$one phase cycle ($\\pm$2$\\pi$ equivalent), i.e., $\\pm$23\\,ps. The average rms is 2 ps, less than one tenth of a phase cycle at 43\\,GHz. The symbols correspond to the following VLBA antennas: B, Brewster; F, Fort Davis; K, Kitt Peak; L, Los Alamos; M, Mauna Kea; N, North Liberty; O, Owens Valley} \\end{figure}  \\noindent Our observations have shown that VLBI differential astrometry at 43\\,GHz provides high-precision relative positions. At this frequency, the astrometric precision is nearly equivalent to the resolution of the maps, and the reference point selected in the source structure might be associated with the core. This makes 43\\,GHz differential astrometry an ideal technique to trace unambiguously the kinematics of the inner regions of the extragalactic radiosources."
    },
    "0002/astro-ph0002302_arXiv.txt": {
        "abstract": "Observations at high spectral and temporal resolution are presented of the dwarf nova EX\\,Dra in outburst. The disk seen in the \\ion{He}{i} line reconstructed by Doppler tomography shows a clear two-armed spiral pattern pointing to spiral shocks in the disk. The Balmer and \\ion{He}{ii} maps also give evidence for the presence of spirals. The eclipse as seen in the red continuum indicates a disk radius of 0.31 times the orbital separation, which might be large enough to explain the observed spiral shocks through excitation by the tidal field of the secondary. The eclipse in the Balmer line profiles, well resolved in our observations, indicates a somewhat smaller disk size (0.25). We discuss the possibility that this is related to an optical depth effect in the lines. ",
        "introduction": "Spiral shocks in accretion disks have been predicted from numerical simulations (Sawada et al. 1986, 1987, R\\'o\\.zyczka \\& Spruit 1993, Yukawa et al. 1997), and analytic considerations (Spruit 1987, Spruit et al. 1987, Larson 1990). They are excited by the tidal field of the secondary if the disk extends far enough into the Roche lobe and can result in two prominent spiral arms. If shock dissipation is the main mechanism damping the wave, it extends over the entire disk and causes accretion at an effective $\\alpha$-value of $0.01(H/r)^{3/2}$ (Spruit 1987, Larson 1990, Godon 1997).  The first observational evidence for shock waves in accretion disks of cataclysmic variables (CVs) was the detection of a clear two-armed structure in the disk of IP\\,Peg during rise to outburst (Steeghs et al. 1997). The spiral pattern, interpreted as evidence for shock waves, has also been seen during outburst maximum (Harlaftis et al. 1999) and early decline of outburst (Morales-Rueda et al. 2000).  At the temperatures expected from dwarf nova disks models, whether in outburst or in quiescence, the predicted spirals are tightly wound and would be hard to detect observationally (Bunk et al. 1990), so that their presence in the observations is somewhat unexpected. Spirals this strong are most naturally explained if the disk temporarily extends rather far into the primary Roche lobe, so that the tidal force of the secondary causes a strong disturbance. A strong non-axisymmetric disturbance, however, would also cause the gas to loose angular momentum quickly (transfered to the secondary), so that the disk would shrink to a smaller radius where the tidal force is weaker. Spirals in disks of CVs would then be understandable if they are a temporary phenomenon, perhaps restricted to outbursts. To test this, more observations of different systems at sufficient spectral resolution and signal-to-noise are needed (Steeghs \\& Stehle 1999).  With high quality spectroscopic studies of different CVs it should also be possible to answer the question if spiral shocks in CV accretion disks are a common phenomenon.  Systems suitable for this purpose would be bright and have frequent outbursts, such as SS\\,Cyg and EX\\,Dra.  EX\\,Dra is a double--eclipsing dwarf nova with a 5-hour orbit (Barwig et al. 1993, Billington et al. 1996, system parameters by Fiedler et al. 1997). There is suggestive evidence for asymmetric structures in the \\ion{He}{i} Doppler map reconstructed from outburst data taken in 1993  (Joergens et al. 2000). From eclipse maps obtained at various stages in the outburst cycle Baptista \\& Catalan (1999) claim that spiral waves form at the early stages of an outburst. We report in this paper on observations at high spectral and temporal resolution during an outburst in 1996. ",
        "conclusions": "We find evidence for spiral structures in the outburst accretion disk of EX\\,Dra similar to those found in IP\\,Peg by Steeghs\\,et\\,al. (1997). The pattern of intensity and velocity perturbations agrees with that predicted from numerical simulations of spiral shock waves (Steeghs\\,\\&\\,Stehle 1999). It is best seen in the \\ion{He}{i} line, somewhat less clearly in the Balmer and \\ion{He}{ii} lines. In EX Dra the pattern appears somewhat less clearly and asymmetric than in IP\\,Peg.  In particular, it is less clear in the \\ion{He}{ii} line, suggesting lower temperatures and shock strengths of the spirals of EX\\,Dra. Possibly the observability of spirals in the \\ion{He}{ii} map is hampered by strong hot spot emission visible in this line.  We derive a disk radius of $r/a=0.31\\pm0.01$ from the  red continuum eclipse light curve. From numerical simulations, Steeghs \\& Stehle find that disk sizes \\mbox{$r_{\\rm d}/a\\,=\\,0.3-0.4$} are needed to excite spirals that are strong enough to generate an observable pattern in the spectra. (Transformation from $r_{\\rm d}/a$ to $r_{\\rm d}/R_{L_1}$ is given by R$_{L_1}=0.53\\,a$ for q=0.75, cp. Plavec \\& Kratochvil 1964.) Since the tidal force is a very steep function of $r/a$, the spirals rapidly become weak at smaller disk sizes. The {\\mk disk size we find here in EX\\,Dra} is at the lower limit of the required size.  Further evidence for the size of the disk in EX\\,Dra in outburst was obtained by Baptista \\& Catalan (1999). Radial intensity distributions presented there show disk radii of 0.30\\,$a$ in quiescence and 0.49\\,$a$ in outburst. The authors see hints of spirals in their eclipse maps, during the early outburst stages. This may be compared with the observations presented here, which show spirals and a disk size of 0.31\\,$a$ three days after the beginning of an outburst.  The eclipses of the Balmer lines in our spectra yield significantly {\\em smaller} disks sizes than the continuum eclipse, $r_{\\rm dB}/a=0.25\\pm 0.02$. Since it is known that the line emission is produced primarily in the outer parts of the disk (e.g. Rutten et al. 1994), one might have expected the disk as seen in the lines to be larger, if anything, than in the continuum.  A possible resolution of this conflict may lie in the optical depth effects affecting the emission lines in systems seen at high inclination. {\\mk The low central intensity of the Balmer lines in high-inclination CVs (often below the continuum) shows that such effects are strong. As shown by Horne \\& Marsh (1986), the effects are strongest for lines of sight parallel and perpendicular to the orbital motion, leading to reduced line emission from these directions compared to intermediate lines of sight (near $45^\\circ$ to the orbit). The bias towards intermediate angles will give the appearance of a somewhat smaller disk size. It is still to be determined if this suggestion also works quantitatively.}"
    },
    "0002/astro-ph0002134_arXiv.txt": {
        "abstract": "Multiwavelength observations of blazars such as Mrk 421 and Mrk 501 show that they exhibit strong short time variabilities in flare-like phenomena. Based on the homogeneous synchrotron self-Compton (SSC) model and assuming that time variability of the emission is initiated by changes in the injection of nonthermal electrons, we perform detailed temporal and spectral studies of a purely cooling plasma system, using parameters appropriate to blazars. One important parameter is the total injected energy ${\\cal E}$ and we show how the synchrotron and Compton components respond as ${\\cal E}$ varies. When the synchrotron and SSC components have comparable peak fluxes, we find that the SSC process contributes strongly to the electron cooling and the whole system is nonlinear, thus simultaneously solving electron and photon kinetic equations is necessary. In the limit of the injection-dominated situation when the cooling timescale is long, we find a unique set of model parameters that are fully constrained by observable quantities. In the limit of cooling-dominated situation, TeV emissions arise mostly from a cooled electron distribution and Compton scattering process is always in the Klein-Nishina regime, which makes the TeV spectrum having a large curvature. Furthermore, even in a single injection event, the multiwavelength light-curves do not necessarily track each other because the electrons that are responsible for those emissions might have quite different lifetimes. We discuss in detail how one could infer important physical parameters using the observed spectra. In particular, we could infer the size of the emission region by looking for exponential decay in the light curves.  We could also test the basic assumption of SSC by measuring the difference in the rate of peak energy changes of synchrotron and SSC peaks. We also show that the trajectory in the photon-index--flux plane evolves clockwise or counter-clockwise depending on the value of ${\\cal E}$ and observed energy bands. ",
        "introduction": "Blazars are a class of flat radio spectrum, core-dominated active galactic nuclei (AGNs). The overall radiation spectra of blazars show two broad peaks in the $\\nu F_{\\nu}$ space; one is between IR and X-rays, and the other in the $\\gamma$-ray regime (e.g., \\cite{vm95}). Flares also have been observed at X- and gamma-ray bands by multiwavelength observations of Mrk 421 (e.g., Macomb et al. 1995; Macomb et al. 1996 for erratum; Buckley et al. 1996) and Mrk 501 (Catanese et al. 1997; Pian et al. 1998). The tremendous luminosity and fast time variabilities from blazars have led to the usual arguments that relativistic motion is occurring in the emitting plasma. Moreover, the favored scenario to explain these sources is that we are viewing nearly along the axis of a relativistically outflowing plasma jet that has been ejected from an accreting super massive black hole (e.g., \\cite{br78}).  Although the origin of these multiwavelength spectra is still under debate, several models on the radiative processes have been put forth, in particular, models of Compton scattering of synchrotron photons or external photons have been developed in recent years (e.g., Bloom \\& Marscher 1996; Inoue \\& Takahara 1996; Ghisellini \\& Madau 1996; Dermer, Sturner, \\& Schlickeiser 1997; Mastichiadis \\& Kirk 1997; Sikora et al. 1997; B\\\"{o}ttcher, Mause, \\& Schlickeiser 1997; Georganopoulos \\& Marscher 1998; Ghisellini et al. 1998). Most of these calculations are either semi-analytic, or for steady state situations, or not including the Compton scattering process self-consistently. The main purpose of this paper is to improve upon this situation. The physics of how energy is dissipated into relativistic particles is, unfortunately, not well understood (see, however, \\cite{rl97}) and will not be treated fully in this paper.  Among various blazar models, synchrotron self-Compton (SSC) models have received a fair amount of attention, by virtue of its simplicity and its possible predictive power. In these models, it is proposed that the nonthermal synchrotron emission forms the radio-through-X-ray continuum, and that the Compton scattering of these (soft) synchrotron photons by the same nonthermal electrons produces the gamma rays ($\\sim$ GeV -- TeVs). In this paper, we focus on the so-called homogeneous SSC model where a spherical blob of uniform relativistic plasma is postulated. Even with such a greatly simplified picture, a number of parameters have to be invoked, whose interplay gives rise to a rich dynamic behavior of the observed radiation. Of particular interest is the correlated variabilities in X-ray and $\\gamma$-ray fluxes, since they represent the tail of nonthermal electrons which have the shortest cooling timescale. Although the generic multiwavelength spectra from radio to TeV can be fitted by a steady state model with fixed parameters (e.g, Kataoka et al. 1999), time-dependent calculations almost always offer stronger constraints. Furthermore, when the self-Compton component contains a comparable or even larger fraction of the radiative energy than the synchrotron component, the whole problem becomes inherently nonlinear and both components need to be calculated simultaneously and self-consistently. This naturally leads to the need of solving coupled, time-dependent, nonlinear particle and photon kinetic-equations. Moreover, by examining the energy-dependence of flare data at gamma-ray energies, one could potentially discriminate between SSC and external Compton-scattering origins of the seed photons (Dermer 1998).  The simplest model for time variability of blazars (Mastichiadis \\& Kirk 1997; hereafter MK97) assumes that electrons obeying a power-law distribution are injected uniformly throughout a relativistically moving blob over an extended period of time, and that electrons cool by both synchrotron radiation and Compton scattering. The blob is assumed not to accelerate or decelerate, and the energy loss by Compton scattering of photons impinging from outside the blob is assumed to be small in comparison with the synchrotron self-Compton loss. MK97 reproduced the qualitative behavior of the energy-dependent lags and the hysteresis diagrams (Takahashi et al. 1996). Much of the work presented here follows closely to the previous study by MK97, but we are using a completely different kinetic code which will be discussed in later sections. Kirk, Rieger, \\& Mastichiadis (1998) further modeled the evolution of synchrotron emission, calculating acceleration and cooling regions separately, though Compton scattering was not included.   In this paper we present a detailed study of the time-evolution of an electron-photon plasma (the positive particles could be either protons or positrons) by solving the kinetic equations numerically. We briefly describe our model in \\S \\ref{sec:codes} and show numerical results in \\S \\ref{sec:results}. Summary is given in \\S \\ref{sec:sum}. ",
        "conclusions": "\\label{sec:sum}  Using a homogeneous synchrotron self-Compton model, we have calculated the time evolution of emission spectra and electron energy distributions when nonthermal electrons are uniformly injected into a relativistically moving plasma blob with constant velocity. We have found that:  (1) When the luminosities of the synchrotron and SSC peaks are comparable, the electron cooling by inverse Compton scattering is not negligible and the system is inherently nonlinear and dynamic. One has to solve the time-dependent, coupled electron and photon kinetic-equations self-consistently. Furthermore, since observations are most sensitive to the peak fluxes of synchrotron and SSC components, accurate treatments of synchrotron emissivity due to the end point effects and inverse Compton scattering in the KN regime are quite essential.  (2) When the cooling time at the maximum particle energy is longer than the  injection timescale ($\\ge R/c$), the light curve of emissions corresponding to the tail of the electron distribution can have short, large amplitude variations but emissions at other wavelengths show considerably longer and smaller amplitude changes. Additionally, spectral evolution is also rather slow. All these features are simply caused by the long cooling time of electrons.  (3) When cooling time at the maximum particle energy is shorter than the dynamic timescale, strong spectral evolutions are observed for both synchrotron and SSC components and short duration flares are obtained in most energy bands.  (4) Generally, the prompt TeV spectrum is curved due to the KN effect and the fact that TeV-production electrons are usually in a cooled distribution. This consideration does not take into account the possible infrared background attenuation of the TeVs, which might cause further curvature in the TeV spectrum. On the other hand, most current TeV observations require an accumulation time probably much longer than the dynamic timescale of the blob, so that it might still be possible to obtain a quasi power-law TeV spectrum by averaging over an evolving spectrum. Further study is needed to address this issue.  (5) We recommend plotting the light curves in a fashion that is logarithmic flux versus linear time. The goal is to look for exponential decays at specific energy bands, which might give direct measurements of the size of the system, as indicated in Figures \\ref{fig:ltcv-e44} -- \\ref{fig:ltcv-e49}.  (6) One has to be cautious about the common belief that light curves in different energy bands should track each other. The electrons responsible for producing specific energy photons might have quite different lifetimes, especially when multiple and closely spaced injections are involved. This complication also applies to the leading/lagging analysis for different photon energy bands.  (7) When high time-resolution spectroscopy is available both in keV and TeV bands, one should be able to prove whether TeV production is via SSC process by comparing the rates of spectral softening as done in Figure \\ref{fig:epkt}.  The primary purpose of this paper is to investigate the radiative signatures in a purely cooling and dynamic system, thus providing a bridge between observations and the detailed but largely unknown physics of particle energization processes. Since we did not address the particle acceleration problem here, in this sense, some of the conclusions drawn above are certainly subject to revisions as our understanding of energy flow in AGNs improves.   In conclusion, we have found that solving time-dependent, coupled electron and photon kinetic-equations provides an easy and efficient way of comparing multiwavelength, time-dependent observations with some simplified SSC models. It has the advantage of naturally combining the spectral and temporal evolutions in a dynamic system, which is very useful when more and more high quality data become available."
    },
    "0002/astro-ph0002120_arXiv.txt": {
        "abstract": "A new determination of the pregalactic helium abundance based on the Magellanic Clouds H~II regions is discussed. This determination amounts to $Y_p = 0.2345 \\pm 0.0030$ and is compared with those derived from giant extragalactic H~II regions in systems with extremely low heavy elements content. It is suggested that the higher primordial value derived by other authors from giant H~II region complexes could be due to two systematic effects: the presence of neutral hydrogen inside the helium Str\\\"omgren sphere and the presence of temperature variations inside the observed volume. ",
        "introduction": "The determination of the pregalactic, or primordial, helium abundance by mass $Y_p$ is paramount for the study of cosmology, the physics of elementary particles, and the chemical evolution of galaxies (e. g. Fields \\& Olive 1998, Izotov et al. 1999, Peimbert \\& Torres-Peimbert 1999, and references therein). In this review we briefly discuss the method used to derive $Y_p$ and its main sources of error as well as a new determination based on observations of the SMC. This determination is compared with those carried out earlier based on extremely metal poor extragalactic H~II regions.  The Magellanic Clouds determination of $Y_p$ can have at least four significant advantages and one disadvantage with respect to those based on distant H~II region complexes: a) no underlying absorption correction for the helium lines is needed because the ionizing stars can be excluded from the observing slit, b) the determination of the helium ionization correction factor can be estimated by observing different lines of sight of a given H~II region, c) the accuracy of the determination can be estimated by comparing the results derived from different points in a given H~II region, d) the electron temperature is generally smaller than those of metal poorer H~II regions reducing the effect of collisional excitation from the metastable 2$^3$ S level of He~I, and e) the disadvantage is that the correction due to the chemical evolution of the SMC is in general larger than for the other systems. ",
        "conclusions": "The $Y_p$ value derived by us is significantly smaller than the value derived by Izotov \\& Thuan (1998) from the $Y$ -- O/H linear regression for a sample of 45 BCGs, and by Izotov et al. (1999) from the average for the two most metal deficient galaxies known (I~Zw~18 and SBS 0335--052), that amount to $0.2443 \\pm 0.0015$ and $0.2452 \\pm 0.0015$ respectively.  The difference could be due to systematic effects in the abundance determinations.  There are two systematic effects not considered by Izotov and collaborators that we did take into account, the presence of H$^0$ inside the He$^+$ region and the use of a lower temperature than that provided by the [O~III] lines. We consider the first effect to be a minor one and the second to be a mayor one but both should be estimated for each object.  From \\ constant \\ density \\ chemicaly \\ homogeneous \\ models \\ computed with CLOUDY we estimate that the maximum temperature that should be used to determine the helium abundance should be 5\\% smaller than $T_e$(O~III). Moreover, if there is additional energy injected to the H~II region $T_e$(He~II) should be even smaller.  Luridiana, Peimbert, \\& Leitherer (1999) produced a detailed photoionized model of NGC~2363. For the slit used by Izotov, Thuan, \\& Lipovetsky (1997) they find an $ICF$(He) = 0.993; moreover they also find that the $T_e$(O~III) predicted by the model is considerably smaller than observed.  {From} the data of Izotov et al. (1997) for NGC~2363, adopting a $T_e$(He~II) 10\\% smaller than $T_e$(O~III) and $\\Delta Y/\\Delta Z = 1.9 \\pm 0.5$ we find that $Y_p = 0.234 \\pm 0.006$.  Similarly, Stasinska \\& Schaerer (1999) produced a detailed model of I~Zw~18 and find that the photoionized model predicts a $T_e$(O~III) value 15\\% smaller than observed, on the other hand their model predicts an $ICF$(He) = 1.00. From the observations of $\\lambda\\lambda$ 5876 and 6678 by Izotov et al. (1999) of I~Zw~18, and adopting a $T_e$(He~II) 10\\% smaller than $T_e$(O~III) we obtain $Y_p = 0.237 \\pm 0.007$; for a $T_e$(He~II) 15\\% smaller than $T_e$(O~III) we obtain $Y_p = 0.234 \\pm 0.007$, both results in good agreement with our determination based on the SMC. Further discussion of these issues is presented elsewhere (Peimbert, Peimbert, \\& Luridiana 2000b).  The primordial helium abundance by mass of $0.2345 \\pm 0.0030 (1 \\sigma)$ --- based on the SMC --- combined with the computations by Copi, Schramm, \\& Turner (1995) for three light neutrino species implies that, at the 95 percent confidence level, $\\Omega_b h^2$ is in the 0.0046 to 0.0103 range.  For $h = 0.65$ the $Y_p$ value corresponds to $0.011 < \\Omega_b < 0.024$, a value considerably smaller than that derived from the pregalactic deuterium abundance, $D_p$, determined by Burles \\& Tytler (1998) that corresponds to $0.040 < \\Omega_b < 0.050$ for $h$ = 0.65, but in very good agreement with the observational estimate of the global budget of baryons by Fukugita, Hogan, \\& Peebles (1998) who find $0.007 < \\Omega_b < 0.038$ for $h$ = 0.65. The discrepancy between $Y_p$ and $D_p$ needs to be studied further.  To increase the accuracy of the $Y_p$ determinations we need observations of very high quality of as many He~I lines as possible to derive $T_e$(He~II), $N_e$(He~II), and $\\tau$(3889) self-consistently. We also need observations with high spatial resolution to estimate the $ICF$(He) along different lines of sight.  \\bigskip  It is a pleasure to acknowledge several fruitful discussions on this subject with: L. Carigi, V. Luridiana, B. E. J. Pagel, M. T. Ruiz, E. Skillman, G. Steigman, S. Torres-Peimbert, and S. Viegas."
    },
    "0002/astro-ph0002316_arXiv.txt": {
        "abstract": "s{ The Alpha Magnetic Spectrometer (AMS) was flown in june 1998 on board of the space shuttle Discovery (flight STS-91) at an altitude ranging between 320 and 390 km. This preliminary version of AMS included an Aerogel Threshold \\cher detector (ATC) to separate $\\bar{p}$ from $e^{-}$ background, for momenta less than 3.5 \\GeVc. In this paper, the design and physical principles of ATC will be discussed briefly, then the performance results of ATC will be presented.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002066_arXiv.txt": {
        "abstract": "We present \\hst\\/ WFPC2 images of Stephan's Quintet which encompass three interacting galaxies and their associated tidal features.  These deep, three-color ($B,V,I$) images indicate recent, massive stellar system formation in various regions within the compact group environment.  We have identified star cluster candidates (SCC) both within the interacting galaxies and in the tidal debris. We compare the SCC colors with stellar population synthesis models in order to constrain cluster ages, and compare the pattern of formation of SCC in different regions to the inferred dynamical history of the group. ",
        "introduction": "The Hickson Compact Groups (HCG; Hickson 1982) are among the densest concentrations of galaxies in the local universe. These high densities combined with relatively low velocity dispersions, $\\sigma\\sim(2-3)\\times10^2$~km~s$^{-1}$ (Hickson et al. 1992), make them active sites of strong galaxy interactions. Interactions are believed to initiate bursts of star cluster formation on many scales from dwarf galaxies along tidal tails to massive star clusters, the progenitors of today's globular clusters. One group in particular, Stephan's Quintet (SQ; also known as HCG~92), is notable for evidence of multiple interactions. \\begin{figure}[t] \\plotfiddle{gallaghers1.eps}{3in}{-90}{50}{50}{-165}{220} \\caption{This $V-$band image is produced from two overlapping pointings of WFPC2. The field-of-view is approximately $3\\farcm7 \\times 2\\farcm5$, and the regions of interest have been labeled.  The fifth member of the group, NGC~7317, is out of the frame to the west. Note that NGC~7320 is a foreground galaxy not relevant to the current discussion.} \\end{figure}  SQ is comprised of five galaxies: NGC~7317, NGC~7318A and B, NGC~7319 and NGC~7320 (see Fig.~1 for galaxy identifications). Based on multiwavelength observations of the group, NGC~7317 and NGC~7320 show no evidence for recent interactions, unlike the other three galaxies (NGC~7320 is a foreground galaxy). In particular, NGC~7318B shows morphological disruption of spiral structure, and a long tidal tail extends from NGC~7319. The interactions have resulted in recent and ongoing star formation as evident from $B-V$ (Schombert et al. 1990), H$\\alpha$ (V\\'{\\i}lchez \\& Iglesias-P\\'aramo 1998) and far-infrared (Xu, Sulentic \\& Tuffs 1999) imaging. Furthermore, in the photometric dwarf galaxy study of Hunsberger, Charlton, \\& Zaritsky (1996), SQ was identified as hosting the richest known system of tidal dwarf galaxy candidates.  From these studies, only the largest star-forming regions were resolved; many of the young stars appeared to be distributed in the diffuse light in the tidal features between the galaxies.  High spatial resolution is required to identify star cluster candidates (SCC) which at the distance of SQ ($z=0.02$; $d\\sim66h^{-1}$~Mpc) are faint point sources on the Wide Field and Planetary Camera 2 (WFPC2). {\\it Hubble Space Telescope} (\\hst) imaging was the obvious next step for investigating the full range in scale of massive star formation structure.  Furthermore, with these images we could investigate whether star clusters form in diverse environments from the inner regions of galaxies to tidal debris tens of kiloparsecs from a galaxy center. ",
        "conclusions": "From \\hst\\ WFPC2 images, we find $\\sim150$ SCC in the environs of SQ.  SCC are found both within the bulges of each of the galaxies NGC~7318A/B and NGC~7319, and also in tidal features. The ages deduced from $B-V$ versus $V-I$ colors of SCC are consistent with the complex interaction scenario outlined by MSM97.  Since only old GCC are found in the centers of NGC~7318A/B, this suggests that recent star formation has not yet occurred there.  Very young SCC are found along the interaction shock front between the ISM of NGC~7318B and the IGM of SQ supporting the hypothesis that this is a recent event.  The spread of ages in SCC found throughout the field is indicative of recurring episodes of interaction-induced star formation."
    },
    "0002/astro-ph0002250_arXiv.txt": {
        "abstract": "Most phenomenological galaxy formation models show a discrepancy between the predicted Tully-Fisher relation and the luminosity function. We show that this is mainly due to overmerging of galaxy haloes, which is inherent in both the Press-Schechter formalism and dissipationless N-body simulations. This overmerging problem be circumvented by including a specific galaxy halo formation recipe into an otherwise standard N-body code. Resolving the overmerging also allows us to include models for chemical evolution and starbursts, which improves the match to observational data {\\it and} renders the modelling more realistic. We use high-redshift clustering data to try and distinguish models which predict similar results at low redshifts for different sets of parameters. ",
        "introduction": "There has been significant recent progress in the study of galaxy formation within a cosmological context, mainly due to a phenomenological approach to this problem. The idea is to start with a structure formation model that describes where and when galactic dark haloes form. A simple description of gas dynamics and star formation provides a means to calculate the amount of stars forming in these haloes. Stellar population synthesis models then provide the spectral evolution, i.e.\\ luminosities and colours, of these galaxies.  Many physical processes are modelled as simple functions of the circular velocity of the galaxy halo. Therefore, the Tully-Fisher relation is the most obvious observational relation to try and predict, as it relates the total luminosity of a galaxy to its halo circular velocity. However, most phenomenological galaxy formation models do not simultaneously fit the I-band Tully-Fisher relation and the B or K band luminosity function. When one sets the model parameters such that the Tully-Fisher relation has the right normalization, the luminosity functions generally overshoot (e.g.\\ Kauffmann, White \\& Guiderdoni 1993; Kauffmann, Colberg, Diaferio \\& White 1999), certainly for the $\\Omega=1$, $H_0=50$ km s$^{-1}$ Mpc$^{-1}$ standard CDM cosmology (in the form given by Davis et al.\\ 1985) that we consider in this paper. Alternatively, when making sure that the luminosity functions matches by changing some of the model parameters, the Tully-Fisher relation ends up significantly shifted with respect to the observed relation (e.g.\\ Cole et al.\\ 1994; Heyl et al.\\ 1995).  In order to keep the modelling as analytical as possible, an extension to the Press \\& Schechter (1974) prescription for the evolution of galaxy haloes (e.g.\\ Bond et al. 1991; Bower 1991; Lacey \\& Cole 1993; Kauffmann \\& White 1993) has been a popular ingredient for implementations of a phenomenological theory of galaxy formation. However, the EPS formalism is designed to identify collapsed systems, irrespective of whether these contain surviving subsystems. This `overmerging' of subhaloes into larger embedding haloes is relevant to the problem of matching both the galaxy luminosity function and the Tully-Fisher relation, as the central galaxy in an overmerged halo is the focus of a much larger cooling gas reservoir than the reservoir that galaxy is to focus of in case its parent subhalo survives. Traditional N-body simulations suffer from a similar overmerging problem (e.g.\\ White 1976), which is of a purely numerical nature, caused by two-body heating in dense environments when the mass resolution is too low (Carlberg 1994; van Kampen 1995).  In order to circumvent these problems, we use an N-body simulation technique that includes a built-in recipe for galaxy halo formation, designed to prevent overmerging (van Kampen 1995, 1997), to generate the halo population and its formation and merger history. This resolves most of the discrepancy sketched above, {\\it and}\\ allows us to make the modelling more realistic by adding chemical evolution and a merger-driven bursting mode of star formation to the modelling. Once stars are formed, we apply the stellar population synthesis models of Jimenez et al.\\ (1998) to follow their evolution. We have enhanced these models with a model for the evolution of the average metallicity of the population, which depends on the starting metallicity. Feedback to the surrounding material means that cooling properties of that material will change with time, affecting the star formation rate, and thus various other properties of the parent galaxy. ",
        "conclusions": ""
    },
    "0002/astro-ph0002470_arXiv.txt": {
        "abstract": "In a spectroscopic follow-up to the VLA FIRST survey, the FIRST Bright Quasar Survey (FBQS) has found 29 radio-selected broad absorption line (BAL) quasars.  This sample provides the first opportunity to study the properties of radio-selected BAL quasars.  Contrary to most previous studies, we establish that a significant population of radio-loud BAL quasars exists.  Radio-selected BAL quasars display compact radio morphologies and possess both steep and flat radio spectra.  Quasars with low-ionization BALs have a color distribution redder than that of the FBQS sample as a whole.  The frequency of BAL quasars in the FBQS is significantly greater, perhaps by as much as factor of two, than that inferred from optically selected samples. The frequency of BAL quasars appears to have a complex dependence on radio-loudness.  The properties of this sample appear inconsistent with simple unified models in which BAL quasars constitute a subset of quasars seen edge-on. ",
        "introduction": "\\label{sectionintro}  Until recently, a search of the astronomical literature would have revealed that broad absorption lines (BAL) are seen in approximately 10\\% of optically selected quasars (\\cite{foltz90}; \\cite{weymann91}) and in exactly 0\\% of radio-loud quasars (\\cite{stocke92}). This dichotomy has puzzled astronomers for years. The BAL quasars can be divided into two classes, high-ionization and low-ionization, which are primarily defined by the presence of broad absorption by C~IV $\\lambda1549$ and Mg~II $\\lambda2800$, respectively. (Note that all low-ionization BAL quasars also show high ionization absorption). The high-ionization BAL (HiBAL) quasars are more common, including 10\\% of all optically selected quasars, while the rarer low-ionization BAL (LoBAL) quasars make up only 1\\% of optically selected quasars.  Prior to the FIRST survey there was only a single example of a LoBAL quasar whose spectrum shows strong absorption by metastable excited states of Fe~II (Q 0059$-$2735, \\cite{hazard87}).  Becker et al.\\ (1997) reported the discovery of two more objects resembling Q0059$-$2735 (FIRST J084044.5+363328 and J155633.8+351758), the second of which is radio-loud. We will refer to these as FeLoBAL quasars.  Both of the new unusual quasars were found by making optical identifications of radio sources from the VLA FIRST survey (Faint Images of the Radio Sky at Twenty-cm, \\cite{becker95}; \\cite{white97}).  Subsequently, Brotherton et al.\\ (1998) identified five more radio-loud BAL quasars (two HiBAL and three LoBAL quasars) from a complete sample of radio-selected ultraviolet excess quasars, firmly establishing the existence of radio-loud BAL quasars.  Lastly, Wills, Brandt, and Laor (1999) have recently suggested that the radio-loud quasar PKS 1004+13 is also a BAL quasar.  For the past five years we have been developing several new radio-selected samples of quasars based on the VLA FIRST survey. The most extensive of these is the FIRST Bright Quasar Survey or FBQS (Gregg et al.\\ 1996, hereafter \\cite{gregg96}; White et al.\\ 2000, hereafter \\cite{white99}).  The goal of the FBQS is to identify all quasars in the FIRST survey brighter than 17.8 on the POSS-I $E$ (red) plate. In the initial 2700 square degrees of the FIRST survey, we defined a sample of 1238 quasar candidates based on positional coincidence between a FIRST source and a POSS-I stellar object (see \\cite{gregg96} and \\cite{white99} for a detailed discussion of the candidate selection criteria).  Spectra have been collected for 90\\% of these candidates, 636 of which have been identified as quasars.  Among these are 29 which display BAL characteristics.  We present the optical spectra and radio spectral indices of these BAL quasars, comparing their radio and optical properties to previous samples of optically selected BAL quasars.  We discuss the selection biases inherent in the survey results and discuss why our sample differs from those based on optically selected samples. ",
        "conclusions": "We have investigated the properties of 29 radio-selected BAL quasars found in the FBQS.  The sample comprises 15 high-ionization BAL quasars, and 14 low-ionization BAL quasars, 4 of which are rare FeLoBALs.  At least 13 are formally radio-loud, unequivocally establishing the existence of a substantial population of radio-loud quasars exhibiting BAL spectral features.  The frequency of BAL quasars appears to be higher than that found in optically selected samples.  Even so, because of selection effects and preferential reddening of LoBAL quasars, the FBQS almost certainly misses additional BAL quasars and the true frequency must be higher. The situation is complicated by indications that the frequency of BAL quasars peaks among the radio-moderate population and decreases for the extremes of radio-loudness.  The BAL quasars show compact radio morphologies, and have a range in radio spectral indices.  The radio properties do not support the popular scenario in which all BAL quasars are normal quasars seen edge-on.  An alternative picture in which BALs are an early stage in the development of new or refueled quasars is preferred."
    },
    "0002/astro-ph0002337_arXiv.txt": {
        "abstract": "The major uncertainties involved in the Chandrasekhar mass models for Type Ia supernovae (SNe Ia) are related to the companion star of their accreting white dwarf progenitor (which determines the accretion rate and consequently the carbon ignition density) and the flame speed after the carbon ignition.  We calculate explosive nucleosynthesis in relatively slow deflagrations with a variety of deflagration speeds and ignition densities to put new constraints on the above key quantities.  The abundance of the Fe-group, in particular of neutron-rich species like $^{48}$Ca, $^{50}$Ti, $^{54}$Cr, $^{54,58}$Fe, and $^{58}$Ni, is highly sensitive to the electron captures taking place in the central layers. The yields obtained from such a slow central deflagration, and from a fast deflagration or delayed detonation in the outer layers, are combined and put to comparison with solar isotopic abundances.  To avoid excessively large ratios of $^{54}$Cr/$^{56}$Fe and $^{50}$Ti/$^{56}$Fe, the central density of the \"average\" white dwarf progenitor at ignition should be as low as \\ltsim 2 \\e9 \\gmc.  To avoid the overproduction of $^{58}$Ni and $^{54}$Fe, either the flame speed should not exceed a few \\% of the sound speed in the central low $Y_e$ layers, or the metallicity of the average progenitors has to be lower than solar.  Such low central densities can be realized by a rapid accretion as fast as $\\dot M$ \\gtsim 1 $\\times$ 10$^{-7}$M$_\\odot$ yr$^{-1}$.  In order to reproduce the solar abundance of $^{48}$Ca, one also needs progenitor systems that undergo ignition at higher densities.  Even the smallest laminar flame speeds after the low-density ignitions would not produce sufficient amount of this isotope.  We also found that the total amount of $^{56}$Ni, the Si-Ca/Fe ratio, and the abundance of some elements like Mn and Cr (originating from incomplete Si-burning), depend on the density of the deflagration-detonation transition in delayed detonations.  Our nucleosynthesis results favor transition densities slightly below 2.2$\\times 10^7$~g cm$^{-3}$. ",
        "introduction": " ",
        "conclusions": "From the very early days of explosive nucleosynthesis calculations, when no direct connection to astrophysical sites was possible yet, it was noticed (Trimble 1975) that the solar Fe-group composition could be reproduced with a superposition of matter from explosive Si burning with about 90\\% originating from a $Y_e$=0.499 source and 10\\% from a $Y_e$=0.46 source.  We have discussed in some detail in the present paper that the central part of SNe Ia could be this second source, while Thielemann et al. (1996) have shown that SN II ejecta with $Y_e$$<$0.498 could cause serious problems.  New results by Hachisu \\etal (1999a, 1999b; see also Li \\& van den Heuvel 1997), which include wind losses in the interaction of binary systems, come to the conclusion that the majority of SNe Ia progenitor systems experience hydrogen accretion on white dwarfs at a rate that has them grow toward the Chandrasekhar mass through steady H- and He-burning. This leads to the single degenerate Chandrasekhar mass scenario (SD/Ch). Binary systems with steady H-burning on accreting white dwarfs and effective accretion rates as high as $\\dot M >$ 1 $\\times$ 10$^{-7}$ M$_\\odot$ yr$^{-1}$ might correspond to observed supersoft X-ray sources. This also would lead to low central ignition densities of the Chandrasekhar mass white dwarf at the thermonuclear runaway ($<$ 2 \\e9 \\gmc), which correspond to the C series of the models discussed in the present paper. A small fraction can deviate from steady H burning and would experience weak hydrogen flashes near the end of the accretion history. Such cases would correspond to our W series of models and even higher ignition densities.  Kobayashi et al. (1998) discussed metallicity effects and their influence on the delay time for the appearance of SNe Ia in galactic evolution.  Our nucleosynthesis results of the present paper are quite consistent with these scenarios and favor case C. The reason is that too high ignition densities lead to a high degree of electron captures and small $Y_e$ values, which would cause an overproduction of $^{54}$Cr and $^{50}$Ti in excess of what is permitted for SNe Ia in a solar mix. We have to make two reservations here. New shell model calculations (Dean et al. 1998; Caurier et al. 1999) indicate that the electron capture rates could be substantially reduced in comparison to the rates by Fuller et al. (1980,1982,1985) employed here. The application of these new capture rates might also permit models of our W series without serious deviations from the allowed central $Y_e$ values.  In addition, such conclusions deal only with the dominant or average SN Ia events. If more neutron-rich nuclei like $^{48}$Ca are also the result of SN Ia nucleosynthesis, an occasional event with a significantly higher ignition density is required, a fact also consistent with the above scenario.  Our nucleosynthesis results imply that for the dominant events the central density of the Chandrasekhar mass white dwarf at thermonuclear runaway must be as low as or lower than \\ltsim 2 \\e9 \\gmc, though the exact constraint depends somewhat on the flame speed. Here is the point where nucleosynthesis predictions can also be of help for providing constraints to the supernova modeling and the burning-front velocity.  A carbon deflagration wave, propagating as slow as $v_{\\rm def}/v_{\\rm s} \\sim 0.015$ or even slightly slower, would be the ideal choice for the neutron-rich species such as $^{54}$Cr and $^{50}$Ti; see cases CS15 and WS30. The latter case makes also clear that slightly larger ignition densities than 1.5 \\e9 \\gmc~(case C) are permitted if the flame speed is increased appropriately.  An acceleration of the flame speed in the outer part of the central layers is permitted to about 3\\% of the sound speed or slightly more, but should clearly stay below 5\\%. For such fast burning fronts the $Y_e$-gradient becomes very flat and too much material in the range $Y_e$ = 0.47-0.485 is produced. This leads to dominant abundances of $^{54}$Fe and $^{58}$Ni, a feature that was already prominent in W7, and causes an excessive overproduction of $^{58}$Ni in galactic evolution.  Our calculations were performed with a constant fraction of the sound speed.  An acceleration of the burning front is expected (Khokhlov 1995; Niemeyer \\& Woosley 1997) and future investigations should include such a time dependence, which is in agreement with the findings discussed above.  Finally, nucleosynthesis can also give clues about the deflagration-detonation transition (DDT). The most obvious consequence of choosing different transition densities is the amount of $^{56}$Ni produced in a SN Ia event.  H\\\"oflich and Khokhlov (1996) find from light curve modeling and spectra that the typical $^{56}$Ni mass should be in the range 0.5-0.7~M$_\\odot$. This agrees with W7. Among the DD models it would ask for a value somewhere between DD1 and DD2 (closer to DD2). Here 3, 2, and 1 stand for DDTs when densities ahead of the flame decrease to 3.0, 2.2, and 1.7 \\e7 \\gmc~. DD1 is excluded for other reasons. The amount of Si-Ca in comparison to Fe is too large in DD1 models in order to compensate for the well-known overproduction of Si-Ca in SNe II during galactic evolution.  Si/Fe ratios in SN Ia models require specific Ia/(II+Ib) ratios in order to obtain a solar mix combined with SN II contributions (see Table 5). DD2 seems to be closest to the present observational constraints for this ratio by Cappelaro et al. (1997).  Small DDT densities favor larger amounts of matter that experience incomplete Si burning.  Low-metallicity constraints require an overproduction of Mn (and Cr) in SNe Ia. These elements are mostly made as $^{55}$Co and $^{52}$Fe (decaying to Mn and Cr), which are favorably produced in incomplete Si burning and would also require a DDT between DD1 and DD2. (One should, however, realize that a fast deflagration could possibly simulate this as well - see the Mn overproduction in Figure 12 - and on the other hand that these numbers would have to be rescaled or reinterpreted in multi-D calculations.) Thus, combining all requirements on the DDT from total Ni yields, Si/Fe and Ia/(II+Ib) ratios, as well as specific elements favored in incomplete Si-burning, we would argue for a DDT density slightly below 2.2 \\e7 \\gmc~, i.e., results between DD1 and DD2. One should, however, be careful with these constraints based on spherically symmetric approximations of the burning front. Full three-dimensional calculations could possibly produce the required ratio of matter from incomplete Si burning and complete Si burning with alpha-rich freeze-out in a different realization.  More extended calculations that make use of the conclusions presented here, including a time dependence of $v_{\\rm def}/v_{\\rm sound}$, the best choice for the DDT density, and a detailed galactic evolution model, replacing our comparisons of only the global yields, will be the next step to undertake. That would also have to include further tests of the metallicity of the exploding object, a topic which has gained in importance with the cosmological interpretation of high-redshift SNe Ia.  For this reason we repeat here in Figure 25a-25c some of the Figures 19-21 with a continuation to smaller mass fractions and purely plotted as a function of velocity, in order to magnify the behavior of the outer layers.  The differences in velocities between CDD and WDD models are negligibly small in these plots.   As mentioned before, the $^{54}$Fe (in the outer layers not affected by electron captures but only by the neutron excess due to the initial metallicity) ranging in velocities up to 15,000 - 19,000 km s$^{-1}$ in the models DD1-DD3, is a strong indicator of metallicity (see Fig. 14).  The minima in S, Ar, and Ca, if observed in spectra, with positions between 16,000 and 21,000 km s$^{-1}$ in the DD1-DD3 models, could give further clues on the deflagration-detonation transition. And finally, any unburned intermediate mass elements at higher velocities would give a clear indication of the metallicity of the accreted matter. Our calculations and plots include only initial compositions of $^{12}$C, $^{16}$O, and $^{22}$Ne. Solar abundances of Ca, Ar, S, Si, or Fe would correspond to mass fractions of $1.2\\times 10^{-4}$, $1.6 \\times 10^{-4}$, $7.6\\times 10^{-4}$, $1.3\\times 10^{-3}$, or $2.3\\times 10^{-3}$.  Thus, any future investigations of very early time spectra, leading to abundance observations at high velocities, would provide strong constraints on the SNe Ia mechanism and the relation to the metallicity of the individual SNe Ia explosion."
    },
    "0002/astro-ph0002101_arXiv.txt": {
        "abstract": "We have conducted observations of the environment around the z=2.15 radio loud quasar 1550-269 in search of distant galaxies associated either with it or the z=2.09 CIV absorber along its line of sight. Such objects will be distinguished by their red colours, R-K$>$4.5. We find five such objects in a 1.5 arcmin$^2$ field around the quasar, with typical K' magnitudes of $\\sim$20.4 and no detected R band emission. We also find a sixth object with K=19.6$\\pm$0.3, and undetected at R, just two arcseconds from the quasar. The nature of all these objects is currently unclear, and will remain so until we have determined their redshifts. We suggest that it is likely that they are associated with either the quasar or the CIV absorber, in which case their properties might be similar to those of the z=2.38 red Ly$\\alpha$ emitting galaxies discovered by Francis et al. (1997). The small separation between the quasar and the brightest of our objects suggests that it may be the galaxy responsible for the CIV metal line absorption system. The closeness to the quasar and the red colour might have precluded similar objects from being uncovered in previous searches for emission from CIV and eg. damped absorbers. ",
        "introduction": "The selection of high redshift galaxies on the basis of their colours has been a growth industry over the last 5 years. Much of this work has concentrated on the selection of high (z$>$3) redshift objects through `dropout' techniques. Such methods have had considerable success (eg. Steidel et al. 1999). However, many of these techniques are reliant on emission in the rest-frame ultraviolet. The UV emission from a galaxy can easily be dominated by a small burst of star formation, or alternatively obscured by a relatively small amount of dust. A population of older quiescent galaxies might thus coexist with the UV selected high redshift objects. Studies of the stellar populations in moderate redshift radio galaxies provide some support for this idea.  A number of authors (eg. Stockton et al. (1995), Spinrad et al. (1997)) have shown that several radio galaxies have ages $>$3--5 Gyr at z$\\sim$1.5, indicating that they must have formed at z$>$5. These results have even been used (Dunlop et al. 1996) to argue that $\\Omega$ must be significantly less than 1.  Old galaxies at moderate redshift, passively evolving from z$>$5 to z=1.5 -- 2.5, would appear as red objects, with R-K' colours $>$4.5. There has been considerable interest in such red objects. Much of this work has centred on red objects found in the fields of known high redshift AGN (eg. Hu \\& Ridgeway (1994), Yamada et al., (1997)). A large survey of the environments of z=1--2 quasars (Hall et al., 1998) finds that such associations are quite common. The present paper attempts to push such studies above z=2. The alternative approach, to study red objects in the field, is also an active area with several surveys dedicated to or capable of finding such objects. See eg. Cohen et al. (1999), or Rigopoulou et al. (in preparation). Red objects need not be old, though. An alternative explanation is that they are heavily obscured, and may contain either a redenned AGN or massive starburst (eg. Dey et al., 1999, Egami et al., 1996). In this context it is interesting to note that several of the objects found in recent deep submm surveys have been identified with very red objects (Smail et al., 1999; Clements et al., in preparation).  Finding emission from the putative galaxies responsible for metal and damped-Ly$\\alpha$ absorption line systems has been the goal of numerous observational programmes. At low redshift there has been considerable success in identifying the galaxies responsible for MgII absorption systems (Bergeron \\& Boisse, 1991; Steidel et al., 1997. At higher redshifts, interest has mostly focussed on the damped-Ly$\\alpha$ absorption systems. Searches for line emission from such objects (eg. Bunker et al. 1999; Wolfe et al. 1992) has met with varying success (Leibundgut \\& Robertson, 1999). Fewer observers have looked in the continuum, but there have been some successes there as well. For example, Aragon-Salamanca et al. (1996) found close companions to 2 out of 10 quasars with damped absorbers in a K band survey. As yet there has been no spectroscopic confirmation of these identifications, but the broad characteristics of these objects, and the small fraction of damped absorbers detected, is consistent with plausible models for the evolution of the galaxies responsible (Mathlin et al., in preparation). Meanwhile, Aragon-Salamanca et al. (1994) looked for counterparts to multiple CIV absorbers lying at z$\\sim$1.6, also using K band observations. They found an excess of K band objects near to the quasars, consistent with their being responsible for the CIV absorption. Once again, there is no spectroscopic confirmation of the assumed redshifts.  The present paper presents the first results of a programme aimed at finding quiescent objects at high redshift (z$\\sim$2--2.5) using optical/IR colour selection techniques. Among the targets observed in an initial test programme was the radio loud quasar 1550-2655, selected as an example radio loud object. The rest of the paper is organised as follows. Section 2 describes our observations, data analysis and presents the results. Section 3 discusses these results and examines three possible origins for the red objects we have found to be associated with 1550-2655. Finally we draw our conclusions.  We assume $\\Omega_M$ = 1, $\\Lambda$=0 and H$_0$=100 kms$^{-1}$Mpc$^{-1}$ throughout this paper. ",
        "conclusions": "At present there are several deficiencies in our data. Firstly we have only obtained limits on the objects R band magnitudes. We must detect them and measure, rather than limit, their R band magnitudes before we can properly determine their colours. Secondly we must obtain spectra for the objects so that we can actually determine, rather than speculate on, their redshift. However, the results presented here suggest that a larger survey of quasar environments, both with and without absorbers, using infrared imagers with adaptive optics correction might shed new light on galaxy populations at large redshift. \\\\~\\\\ {\\bf Acknowledgments} This paper is based on observations made at the European Southern Observatory, Chile. It is a pleasure to thank Nick Devillard for his excellent Eclipse data reduction pipeline, and E. Bertin for SExtractor. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. I would like to thank Amanda Baker and Garry Mathlin for useful discussions, and the anonymous referee for helpful comments on an earlier version. This work was supported in part by an ESO fellowship, EU TMR Network programme FMRX-CT96-0068 and by a PPARC postdoctoral grant."
    },
    "0002/astro-ph0002271_arXiv.txt": {
        "abstract": "\\nop We estimate the contributions to the cosmic microwave background radiation (CMBR) power spectrum from the static and kinematic Sunyaev-Zel'dovich (SZ) effects, and from the moving cluster of galaxies (MCG) effect. We conclude, in agreement with other studies, that at sufficiently small scales secondary fluctuations caused by clusters provide important contributions to the CMBR. At $\\ell \\gtrsim 3000$, these secondary fluctuations become important relative to lensed primordial fluctuations. Gravitational lensing at small angular scales has been proposed as a way to break the ``geometric degeneracy'' in determining fundamental cosmological parameters. We show that this method requires the separation of the static SZ effect, but the kinematic SZ effect and the MCG effect are less important. The power spectrum of secondary fluctuations caused by clusters of galaxies, if separated from the spectrum of lensed primordial fluctuations, might provide an independent constraint on several important cosmological parameters. ",
        "introduction": "\\label{S:INTRO}   The power spectrum of the cosmic microwave radiation (CMBR) carries much cosmological information about primordial density fluctuations in the early Universe. As photons leave the last scattering surface and travel across the Universe, however, these brightness fluctuations are modified by intervening structures, causing secondary fluctuations, which we would expect to become more important on small angular scales.  The power spectrum of the CMBR alone can be used to determine cosmological parameters. Recently it has been shown, however, that a geometrical degeneracy effect prevents some combinations of cosmological parameters from being disentangled by the power spectrum alone (\\cite{BondEfTeg97}; \\cite{EfstathBond98}; \\cite{ZaldarrSperSel97}). The primordial density fluctuations and matter content determine the positions and magnitudes of the Doppler peaks at the last scattering surface. These fluctuations are transferred to apparent angular scales determined by their angular diameter distance. As a result, cold dark matter (CDM) models with the same primordial density fluctuations, matter content, and angular diameter distance can not be distinguished. The models are ``effectively degenerate'' in the sense that their power spectrum is degenerate for parameter determination on intermediate and small scales. Although the observed power spectrum also depends on the time variation of the metric, via the integrated Sachs-Wolfe effect, this breaks the degeneracy only at large angular scales. Unfortunately observations of the power spectrum do not provide strong constraints on models at large scales, since this is where the statistics of the data are dominated by the cosmic variance due to the fact that we have only one realization of our cosmological model, the Universe itself, and we encounter a sampling problem). Therefore we can determine, for example, only combinations such as $\\Omega_0 h^2$ and $\\Omega_b h^2$ (where $\\Omega_0$ and $\\Omega_b$ are the $z = 0$ matter and baryon density parameters and $h$ is the dimensionless Hubble parameter).  It has been noticed that gravitational lensing can break the geometric degeneracy at small angular scales, $\\ell \\gtrsim 2000$, in such a way that the cosmological parameters can be determined separately (\\cite{MetcalfSilk98}; \\cite{StomporEfstath98}). The effect of gravitational lensing on the CMBR was studied by several authors (see for example \\cite{BlanchardSchneider87}; \\cite{MetcalfSilk97}; \\cite{Seljak97}). Static gravitational lenses do not change a smooth CMBR, but the fluctuations get distorted by lensing. As a result, power from the acoustic peaks is transferred to small angular scales, conserving the variance of the spectrum. The amount of power transferred depends on the cosmological model, thus, in principle, we can determine separately $\\Omega_0$, $\\Omega_b$, and $h$. This method makes use of the small angular scale part of the power spectrum, where the amplitude of primordial fluctuations is declining and secondary fluctuations are becoming more important. The question naturally arises: How do contributions to the power spectrum from secondary fluctuations influence parameter determination based on the small scale CMBR power spectrum?   The most important secondary fluctuations are the thermal static and kinematic Sunyaev-Zel'dovich (SSZ and KSZ) effects (\\cite{sz80}), the Rees-Sciama (RS) effect (\\cite{ReesSciama68}), the moving cluster of galaxies (MCG) effect (\\cite{BirkGull83}; \\cite{GurvitsMitrof86} ; \\cite{PyneBirkinshaw93}), point sources (\\cite{ToffZAB99}), and, if the Universe was re-ionized at some early stage, the Ostriker-Vishniac effect (\\cite{OstrVish86}; \\cite{Vish87}). In this paper we concentrate on secondary effects caused by clusters of galaxies. Since detailed reviews are available on the SZ and the MCG effects (\\cite{Reph95}; \\cite{Birk98}), here we just summarize their major features. The thermal SZ effect is a change in the CMBR via inverse Compton scattering by electrons in the hot atmosphere of an intervening cluster of galaxies. We use the terms kinematic or static SZ effect depending on whether or not the intracluster gas possesses bulk motion. To date only the static thermal SZ effect has been detected (\\cite{Birk98}).  The MCG effect is a special type of RS effect, due to the time-varying gravitational field of a cluster of galaxies as it moves relative to the rest frame of the CMBR. Unlike the original RS effect, the MCG effect is not caused by intrinsic variation of the gravitational field, so that in the rest frame of the cluster, the photons fall into and climb out of the same gravitational field. However, in the rest frame of the cluster the CMBR is not isotropic, but has a dipole pattern, being brighter in the direction of the cluster peculiar velocity vector. Photons passing the cluster are deflected towards its center. Thus in the direction of the cluster peculiar velocity vector (ahead of the cluster) one can see a cooler part of the dipole pattern. Towards the tail of the cluster, one can see a brighter part of the dipole (ahead of the cluster). When transferring back to the rest frame of the CMBR, we transfer the dipole out, but the fluctuations remain, showing a bipolar pattern of positive and negative peaks. At cluster center there is no deflection, thus there is no effect. The amplitude of the MCG effect is proportional to the product of the gravitational deflection angle and the peculiar velocity of the cluster.  The most important characteristics of the SSZ, KSZ, and MCG effects in the context of cosmology is that their amplitudes do not depend on the redshift of the clusters causing the effect. Using thermodynamic temperature units, their maximum amplitude are about 500 $\\mu$K, 20 $\\mu$K, and 10 $\\mu$K, respectively. The SSZ and KSZ effects have the same spatial dependence as the line of sight optical depth, the MCG effect has a unique bipolar pattern. Assuming a King approximation for the total mass and an isothermal beta model for the intracluster gas, the full width at half maximum (FWHM) of the SSZ and KSZ effects $\\approx (2 - 4)\\,r_c$, where $r_c$ is the core radius, depending on $\\beta$, which is typically between 2/3 and 1. The MCG effect has a much larger spatial extent, with FWHM for each part of the bipolar distribution $\\approx 10\\, r_c$.  The spectra of the effects are also important: the SSZ effect has a unique spectrum which changes sign from negative to positive at about 218 GHz. The KSZ and MCG effects have the same frequency dependence as the primordial fluctuations. The most important difference between the SZ effect and the MCG effect is that the SZ effect is caused by intracluster gas, the MCG effect is caused by gravitational lensing by the total mass regardless the physical nature of that mass. Therefore the SZ effect only arises from clusters with intracluster gas. Clusters can produce significant MCG effects even if devoid of intracluster gas.   The effects of clusters of galaxies on the CMBR in a given cosmology have been a subject of intensive research since the late 1980s. There are several different ways of extracting information from these effects. Source counts of the SSZ effect were estimated by using the Press-Schechter mass function (PSMF) and scaling relations (\\cite{ColeKaiser88}; \\cite{Markevitchet92}; 1994; \\cite{MakiSuto93}; \\cite{BartSilk94}, De Luca, Desert \\& Puget 1995, \\cite{ColaMRV94}; 1997, \\cite{sutoet99}). The importance of the SSZ effect was demonstrated and it was shown that thousands of detections are expected with the next generation of satellites. Contributions to the CMBR from the RS and the KSZ effects were derived by Tuluie, Laguna \\& Anninos (1996) and Seljak (1996) for CDM models with zero cosmological constant. Tuluie et al. used N-body simulations and a ray-tracing technique, Seljak used N-body simulations and second order perturbation theory. Contributions from the SSZ and KSZ effects originating from large scale mass concentrations (superclusters) were studied by Persi et al. (1995). Bersanelli et al. (1996), in their extensive study of the CMBR for the Planck mission, estimated the contribution to the power spectrum from the SSZ and KSZ effects. Aghanim et al. (1998) estimated the effects of the KSZ and MCG effects on the CMBR including their contributions to the CMBR power spectrum. Aghanim et al. simulated $12.5^\\circ \\times 12.5^\\circ$ maps with pixel size of $1.5^\\prime \\times 1.5^\\prime$. They used the PSMF normalized to X-ray data (assuming an X-ray luminosity-mass relation). The total mass was assumed to have a Navarro-Frenk-White profile (\\cite{Navarroet97}), and the intracluster gas was assumed to follow an isothermal beta model distribution. The time evolution of the electron temperature and the core radius were assumed to follow models of Bartlett \\& Silk (1994), which are based on self-similar models of Kaiser (1986). According to Aghanim et al. (1998)'s results, the KSZ effect is many orders of magnitude stronger than the primordial CMBR on small angular scales, and therefore the effect would prevent the use of the power spectrum to break the geometric degeneracy. Atrio-Barandela \\& Mucket (1998) estimated the power spectra of the SSZ effect in a standard dark matter dominated model with different lower mass cut-offs. Contributions to the power from the Ostriker-Vishniac effect in CDM models were estimated by Jaffe \\& Kamionkowski (1998).   In this paper we estimate the contributions to the CMBR power spectrum from the SSZ, KSZ, and MCG effects on small angular scales adopting cold dark matter dominated models. In our models we assumed scale invariant primordial fluctuations with a processed spectrum having a power law form on cluster scales with a power law index of $n_P = -1.4$ (Bahcall \\& Fan 1998). This maybe used as a first approximation as long as contributions from very low and/or very high mass clusters are small (cf. our discussion about mass cut offs at section~\\ref{s:PSMF}). We use three representative models in our study: {\\it Model 1}, open CDM (OCDM) model: a low density open model with $\\Omega_0 = 0.2$, $\\Lambda = 0$, $\\sigma_8 = 1.2$; {\\it Model 2}, flat lambda CDM (\\L CDM) model: a low density flat model with $\\Omega_0 = 0.2$, $\\Lambda = 0.8$, $\\sigma_8 = 1.35$; {\\it Model 3}, standard CDM (SCDM) model: a flat model with $\\Omega_0 = 1$, $\\Lambda = 0$, and $\\sigma_8 = 0.65$.  In Section~\\ref{s:method} we outline our method of estimating the power spectra of secondary fluctuations caused by clusters of galaxies and discuss our normalization method for the PSMF. Sections~\\ref{s:PSMF} and ~\\ref{s:physpar} describe how we used the PSMF and the scaling relations to obtain masses and other physical parameters of clusters. In section~\\ref{s:PowerSpectr} we present the spherical harmonic expansion of the SSZ, KSZ and MCG effects, and our method of estimating their power spectra. Section~\\ref{s:Simulation} describes our simulations to evaluate the integrals over clusters. Sections~\\ref{s:Results} and \\ref{s:Discussion} present our results and discuss the differences from previous work. ",
        "conclusions": "\\label{s:Discussion}   An observed power spectrum is made up from the sum of all astrophysical effects and noise. We rely on the different frequency and/or power spectra of the secondary effects to separate these foregrounds from the primordial CMBR signal (\\cite{Tegmark98}).  Of the effects discussed here, it should be easy to separate the SSZ effect by using multi-frequency measurements of its unique spectrum. The separation of primordial fluctuations in the CMBR and fluctuations caused by the KSZ and MCG effects is more difficult since their frequency spectra are the same. Optimal filters have been designed to separate the KSZ effect (\\cite{HaehneltTegmark96}; \\cite{Aghanimet97}): here it helps to know the SSZ effect for the same cluster, since that would give us a position and even an estimate for the expected amplitude of the effect. Aghanim et al. (1998) discussed methods to separate the MCG effect: this is facilitated by its unique dipole pattern with sharp peaks (Figure~\\ref{F:MCG1}). Primordial fluctuations are usually assumed to be gaussian, where the probability of getting such a strongly peaked bipolar pattern is small, and we would expect the strong small angular scale gradient near a known cluster of galaxies to be a definite indication of the presence of the MCG effect. Also, knowing the position of clusters helps to find the effect. However, contributions from other effects, such as early ionization and discrete radio sources causes further confusion, and may be expected to make it difficult to determine the power from the SZ or MCG effects.   We analyzed the contributions to the power spectrum from the SSZ, KSZ and MCG effects to check their impact on the determination of cosmological parameters, especially at large $\\ell$ where gravitational lensing may break the geometric degeneracy. In Figure~\\ref{f:Lens01} we show the small scale lensed primordial fluctuation power spectra of our three models (OCDM, SCDM, \\L CDM; solid lines) with power spectra resulting from the sum of fluctuations due to the lensed primordial CMBR and the SSZ effect (long dashed lines), and from the sum of the lensed primordial CMBR, the KSZ and the MCG effects (short dashed lines). According to our models, if the fluctuations due to the SSZ effect are fully separated, the KSZ and MCG effects do not prevent the use of this part of the power spectrum to break the geometric degeneracy and distinguish between different CDM models. Note that normalization at the first Doppler peak, rather than the usual {\\it COBE} normalization, would lower the contributions to the power spectrum from primordial fluctuations in a \\L CDM model relative to those from a SCDM model, and thus secondary effects would become more important relative to the primordial CMBR fluctuations.  Our simulations also show that the power spectrum of the SSZ effect may itself be used to break the geometric degeneracy. Since the separation of the SSZ effect from other secondary effects should be straightforward, we should be able to determine the power spectrum of the SSZ effect alone. As can be seen from Figure~\\ref{f:PowerSp3}, this power spectrum depends on $\\Omega_0 h^2$ and $\\Omega_b h^2$, providing an additional constraint on these parameters. We note, however, that the amplitude of the SSZ effect is model dependent. Since the contributions to the power spectrum from the SZ and MCG effects are model dependent, to evaluate fully their power spectra we need a better observationally-supported model for the intracluster gas. Sensitive, high-resolution all-sky, X-ray observations could map the emission from intracluster gas up to high redshift providing strong constraints on gas formation and evolution and thus a good basis for modeling the SSZ and KSZ effects (\\cite{Jahoda.et97}). Number counts of clusters based on their SSZ effect can also be used to constrain cosmological models (\\cite{sutoet99}).  There are many possibilities of using observations to break the geometric degeneracy. For example measurements of the CMBR polarization, the Hubble constant, or light curves of Type Ia supernovae have been discussed (Zaldarriaga et al. 1997; \\cite{EisensteinHuTeg98}; \\cite{Tegmarket98}). Also, combination of measurements of the SSZ effect and thermal bremsstrahlung (X-ray) emission from clusters can be used to determine the Hubble constant for a large number of clusters, providing a statistical sample which might enable us to determine the Hubble constant, and perhaps the acceleration parameter, to good accuracy (\\cite{Birk98}).   Secondary fluctuations introduce non-gaussianity into the primordial spectrum at small scales. This non-gaussianity should be taken into account when estimating CMBR non-gaussianity at these scales. Winitzky (1998) estimated the effect of lensing and concluded that Planck may observe non-gaussianity due to lensing near the angular scale of maximum effect, $\\sim 10^\\prime$. Other processes, including the SSZ, KSZ, and especially the MCG effect, introduce a highly non-gaussian signal as is easily seen for the MCG effect on Figure~\\ref{F:MCG1}. A similar non-gaussian pattern arises from moving cosmic strings (the Kaiser-Stebbins effect, compare our Figure~\\ref{F:MCG1} to Figure 6a of \\cite{MagueijoLewin97}). Our results indicate that at $\\ell \\gtrsim 10^4$ the MCG effect might be comparable in strength to the primordial fluctuations. Evidence for non-gaussianity has been reported by Ferreira, Magueijo \\& Gorski (1998) and Gaztanaga, Fosalba, \\& Elizalde (1997) at angular scales $\\ell \\approx 16$ and $\\ell \\approx 150$. They do not exclude the possibility that this non-gaussianity has been introduced by foregrounds, but our results show that clusters can not introduce detectable non-gaussianity on such scales (Figure~\\ref{f:PowerSp3}).   We convolved our theoretical results (Figures~\\ref{f:PowerSpOCDM} - \\ref{f:PowerSpSCDM}) with the expected point spread functions (PSFs) of instruments on the MAP and Planck missions to estimate the level of the secondary fluctuations caused by clusters of galaxies on the observable power spectrum. The observed $C_{\\ell}$ values become  \\begin{equation} C_{\\ell}^{obs} = C_{\\ell} W_{\\ell} , \\end{equation} where the $W_{\\ell}$ values are the Legendre coefficients of the PSF. For an assumed gaussian response function and in the small angle approximation, the $W_{\\ell}$ coefficients are  \\begin{equation} W_{\\ell} = e^{- \\sigma^2 (\\ell + 1/2)^2} , \\end{equation} where $\\sigma = h / 2 \\sqrt{ \\ln 4}$, and $h$ is the FWHM of the beam (for a detailed description of window functions and $W_{\\ell}$ coefficients, see White \\& Srednicki 1995). The observed rms fluctuations then become  \\begin{equation} \\label{E:DT_rms} \\langle \\Delta T/ T_0 \\rangle_{rms}^2 = \\sum_\\ell { 2 \\ell + 1 \\over 4 \\pi}  C_\\ell W_{\\ell} . \\end{equation}  In general, contributions from unresolved cluster static effects add to provide a cumulative contribution to the CMBR power spectrum. Contributions from the KSZ and MCG effects from unresolved sources tend to cancel. In the case of the MCG effect this is because each unresolved source contribution would be zero owing to the dipole spatial pattern of the effect. For small-scale KSZ effects there are several sources in the field of view of the telescope, and different sources have positive or negative contributions depending on the sign of their line of sight peculiar velocity, and therefore they tend to cancel each other. The larger the beam size, the more effective is the cancellation of the MCG and KSZ effects. Note that the spatial extension of the MCG effect is much larger than that of the KSZ effect, so many clusters may be unresolved in their KSZ and resolved in their MCG effect. The MCG effect might be relatively more important at high redshifts, since it does not require a well-developed cluster atmosphere.  In Figures~\\ref{f:MAP94} and \\ref{f:PLANCK353} we show the contributions to the power spectrum from primordial fluctuations, and the SSZ, KSZ and MCG effects, convolved with the PSF of the planned $\\nu = 94$ GHz receiver on MAP, and the planned $\\nu = 353$ GHz bolometer on Planck. The amplitude of the fluctuations from the the SSZ effect is negative at $\\nu = 94$ GHz and positive at $\\nu = 353$ GHz, but only the the absolute value of the effect contributes to the power spectrum. The different maximum $\\ell$ values for the MAP and Planck systems ($\\ell_{max} \\sim$ 1000 and 2000, respectively) can clearly be seen on Figures~\\ref{f:MAP94} and \\ref{f:PLANCK353}. Because of these cutoffs, the observable power spectrum is dominated by primordial fluctuations at all $\\ell$ for these missions. According to our results, the SSZ effect may cause a 1\\% enhancement in the amplitude of the Doppler peaks, which is at the limit of the sensitivity of the MAP and Planck missions. From the analysis of the power spectrum, this would lead to an overestimation of the parameter $\\Omega_0 h^2$ by about 1\\%,. The shift in the position of peaks as a function of $\\ell$ caused by the SSZ effect is less important since the spectrum of the SSZ effect has only a weak dependence on $\\ell$.  In Table~\\ref{T:RMS_ALL} we show the $(\\Delta T/T)_{rms}$ values of the contributions to the CMBR from the SSZ, KSZ, and MCG effects convolved with the the 94 GHz MAP and 353 GHz Planck receivers for our three models (OCDM; \\L CDM; SCDM). As a comparison, we display the corresponding rms values of the primordial fluctuations. The rms values of all these secondary effects are an order of magnitude smaller than rms values from primordial fluctuations. The most important contribution is that of the SSZ effect at these frequencies. The KSZ and MCG effects give similar contributions with the KSZ effect being about a factor of two stronger. Aghanim et al. (1998)'s results for the rms values of the MCG effect is a factor of 10 (SCDM) or a factor of 3 (OCDM and \\L CDM) larger than our results. Note, however, that rms values give only a crude estimate of the magnitude of the effects. At large angular scales the primordial fluctuations are about 100 (for the SSZ effect) or $10^5 - 10^7$ (KSZ, MCG effects) times stronger than the secondary fluctuations.  An ideal observation to measure the contribution to the power spectrum from the SSZ effect would use high angular resolution ($\\ell \\gtrsim 7000$) and high frequency ($\\nu \\gtrsim 250$ GHz). SuZIE probes the power spectrum at angular scale $\\ell \\approx 7500$ at 140 GHz. The 2$\\sigma$ upper limit on the power at this scale from SuZIE is $\\ell (\\ell+1) C_\\ell \\leq  1.4 \\times 10^{-9}$ (\\cite{Gangaet97}). Unfortunately our models suggest that at this frequency the primordial contribution to the power spectrum is about 10 times stronger than that from the SSZ effect. A promising experiment is SCUBA, which probes the anisotropies at angular scale $\\ell \\approx 10000$ and frequency 348.4 GHz. Their preliminary 2$\\sigma$ upper limit on the power spectrum at this scale is $\\ell (\\ell+1) C_\\ell \\leq  4.7 \\times 10^{-8}$. Much further work is planned, and should lower this limit by a factor of 3-10 (\\cite{BorysCS98})."
    },
    "0002/astro-ph0002047_arXiv.txt": {
        "abstract": "We present results from photoionization models of low-metallicity \\h regions. These nebulae form the basis for measuring the primordial helium abundance. Our models show that the helium ionization correction factor (ICF) can be non-negligible for nebulae excited by stars with effective temperatures larger than 40,000~K. Furthermore, we find that when the effective temperature rises to above 45,000~K, the ICF can be significantly negative. This result is independent of the choice of stellar atmosphere. However, if an \\h region has an \\3/\\1 ratio greater than 300, then our models show that, regardless of its metallicity, it will have a negligibly small ICF. A similar, but metallicity dependent, result was found using the \\3/H$\\beta$ ratio. These two results can be used as selection criteria to remove nebulae with potentially non-negligible ICFs. Using our metallicity independent criterion on the data of \\citet{izo98} results in a 20\\% reduction of the rms scatter about the best fit $Y-Z$ line. A fit to the selected data results in a slight increase of the value of the primordial helium abundance. ",
        "introduction": "\\label{sec:intro}  An accurate measurement of the primordial helium abundance would be an important test of standard big bang nucleosynthesis \\citep{oli97}, and would also constrain the values of the photon-to-baryon ratio and $\\Omega_b$ \\citep{oli99}. The traditional procedure to measure the primordial helium abundance is to make use of the correlation between the helium mass fraction ($Y$) and metal abundance ($Z$). This correlation is then extrapolated to zero metallicity to estimate the primordial mass fraction of helium, $Y_p$. Spectroscopic observations of bright, low-metallicity extragalactic \\h regions provide the data for these studies \\cite[e.g.,][]{oli95,oli97, izo94,izo97,izo98,tor89,skil98}.  To be cosmologically useful the value of $Y_p$ has to be determined to better than 5\\%. Fortunately, abundance determinations from measurements of line ratios is relatively straightforward \\citep{pei75,ben99}, and can, {\\it in theory}, give the desired accuracy. However, to reach the needed level of precision, any systematic errors involved with target selection, observations, and data analysis must be identified and corrected. Many such systematic errors have already been identified \\citep{dav85,din86,pag92,skil94,pei96,izo97,ste97,skil98}, but any errors resulting from the so-called ``ionization correction factor'' (ICF) have so far been assumed to be small. The ICF corrects for the fact that some amount of atomic (i.e., unseen) helium might be present in ionized regions of hydrogen \\citep{ost89,pei75}. This correction traditionally has been assumed to be zero because measurements of the primordial helium abundance employ observations of bright extragalactic \\h regions. These regions are excited by clusters of young stars with effective temperatures greater than 40,000~K. Calculations by \\citet{sta90} and \\citet{pag92} showed that the helium ICF should be negligibly small for these \\h regions. As a result, recent determinations of $Y_p$ have assumed that the helium ICF is small.  Very recently, \\citet{arm99} presented calculations that showed that \\h regions excited by stars with temperatures greater than 40,000~K can have non-negligible ICFs. \\citet{arm99} found that the ICFs were often negative (i.e., the helium ionized zone is {\\bf larger} than the hydrogen one; \\citet{sta80,sta82}, \\citet{pen86}) for the hardest stellar continua. These results were confirmed by \\citet{vie99}. In this paper, we follow up on the work of \\citet{arm99}, and develop observational diagnostics of when the He ICF is important and when it can be ignored. We then apply these diagnostics to the data of \\cite{izo98} to illustrate how our technique can improve the precision of the measurement of $Y_p$.  We describe our calculations in \\S~\\ref{sec:calc}, and our results in \\S~\\ref{sec:results}. The main results are summarized in \\S~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl} In this paper we have shown the following: \\begin{enumerate}  \\item There can be a non-negligible ICF correction for \\h regions excited by stars with temperatures greater than 40,000~K. At temperatures higher than 45,000~K, the ICF is preferentially negative. This result is independent of the atmosphere of the O star.  \\item There is a simple procedure to determine if an ICF correction needs to be made for a given \\h region. If the \\3/\\1 ratio is greater than 300, then no correction is needed. This criterion is independent of metallicity. If the \\1 line cannot be measured, then there is a metallicity dependent cutoff (Eq.~\\ref{eq:zcutoff}) that can be used with the \\3 line.  \\item Applying the metallicity independent criterion to the data of \\citet{izo98} results in reducing the rms scatter about the best fit $Y-Z$ line by 20\\%. This will help remove systematic errors relating to unrecognized ICF effects, and ought to improve the reliability of the $Y_p$ determination. Furthermore, an analysis of the selected data gives a larger value of $Y_p$ than was originally measured, which is closer to the theoretically expected value. \\end{enumerate}"
    },
    "0002/astro-ph0002517_arXiv.txt": {
        "abstract": "We report the detection of a broad 22 $\\mu$m emission feature in the Carina nebula \\ion{H}{2} region by the Infrared Space Observatory (ISO) Short Wavelength Spectrometer. The feature shape is similar to that of the 22 $\\mu$m emission feature of newly synthesized dust observed in the Cassiopeia A supernova remnant. This finding suggests that both of the features are arising from the same carrier, and that supernovae are probably the dominant production source of this new interstellar grain. A similar broad emission dust feature is also found in the spectra of two starburst galaxies from the ISO archival data. This new dust grain could be an abundant component of interstellar grains and can be used to trace the supernova rate or star formation rate in external galaxies. The existence of the broad 22 $\\mu$m emission feature complicates the dust model for starburst galaxies and must be taken into account correctly in the derivation of dust color temperature. Mg protosilicate has been suggested as the carrier of the 22 $\\mu$m emission dust feature observed in Cassiopeia A. The present results provide useful information in studies on chemical composition and emission mechanism of the carrier. ",
        "introduction": "Supernovae have been suggested besides evolved stars as one of the major sources of interstellar dust (see Gehrz 1989, Jones and Tielens 1994, Dwek 1998 for review). Supporting evidence includes observations of dust condensation in the ejecta of SN 1987A (Moseley et al. 1989, Whitelock et al. 1989, Dwek et al. 1992, Wooden et al. 1993), and those of the newly synthesized dust in the Cassiopeia A (Cas A) supernova remnant (Arendt, Dwek, \\& Moseley 1999). The dust formation mechanism and the amount of dust that is formed in supernovae are still poorly known. Observations of SN 1987A and Cas A showed that the mass of the newly formed dust is much less than expected, and the discrepancy may be due to the fact that most of the dust is cold and cannot be detected in the far-infrared (Dwek 1998, Arendt et al. 1999). Finding an abundant dust component in the interstellar medium (ISM) which is formed only in supernovae will support the hypothesis that supernovae are a major source of interstellar dust. Furthermore, since the amount of this specific grain is proportional to the number of supernova, its total mass in the ISM can be used as a tracer of the supernova rate or star formation rate in external galaxies. In this Letter we report the detection of a broad 22 $\\mu$m emission dust feature in the Carina nebula \\ion{H}{2} region by the  ISO guaranteed time observations. We found that the shape of the present 22 $\\mu$m emission dust feature is similar to the 22 $\\mu$m emission feature observed in Cas A. We also found a similar emission feature in two starburst galaxies from the ISO archival data. ",
        "conclusions": "Evolved stars and supernovae have been suggested as the major production sources of interstellar dust. Past observations of evolved stars have found a number of dust features in the near to far-infrared ranges (see Waters et al. 1999 for a recent review). However, the broad 22 $\\mu$m emission feature that we found in Carina nebula \\ion{H}{2} region has never been reported in evolved stars. On the other hand, the present broad 22 $\\mu$m emission feature is quite similar to the emission feature of newly synthesized dust observed in Cas A, suggesting that both of these features arise from the same dust grain, and that supernovae are probably the major production source of this new interstellar grain. The non-detection of the 22 $\\mu$m feature in SN 1987A (Moseley et al. 1989) does not make the latter suggestion less convincing, since the infrared emission in SN 1987A probably arises from optically thick clumps. Lucy et al. (1991) and Wooden et al. (1993) suggest that the infrared emission in SN 1987A is dominated by the dust in the optically thick clumps, and the low density small grains in the interclump medium contribute to the visual extinction. With this model, the infrared emission in SN 1987A is a graybody emission, but the visual extinction is not.  We would expect to find the 22 $\\mu$m dust feature in astronomical sources with high supernova rate if supernovae are the major production source of this new interstellar grain.  Starburst galaxies are an ideal place to search for. From the ISO archival data we found that two starburst galaxies, M82 and NGC7582, show a similar 22 $\\mu$m emission feature. Figure 4 shows the SWS spectrum of the nuclear region of NGC7582, a narrow-line  X-ray galaxy with strong starburst in the central kpc (Radovich et al. 1999, and references therein). The 20 to 30 $\\mu$m emission is mostly or completely arising from the broad 22 $\\mu$m emission feature. The spectrum of NGC7582 was taken by the SWS AOT01 with the speed of 2. We processed the data through the OLP 8.4 and reduced with the SWS IA package in a way similar to the Carina nebula spectra. The feature intensity in M82 (not shown) is much weaker, about 10$\\%$ of the 18 -- 28 $\\mu$m emission if the continuum is assumed to pass through the 18 and 28 $\\mu$m data points. Two other starburst galaxies, NGC253 and Circinus may also have a 22 $\\mu$m feature, but they are further weak in intensity and more observations are needed to confirm it.  The findings of the 22 $\\mu$m dust feature in \\ion{H}{2} regions and starburst galaxies suggest that this new grain could be an abundant component of interstellar dust. If the amount of this interstellar grain in the ISM is supposed to be proportional to the number of supernovae, its total mass in the ISM can be used as a tracer of the supernova rate or star formation rate in external galaxies. Studies of a large sample of starburst galaxies are required to confirm the above relationship. Only a limited number of galaxies have been observed by the SWS full grating scan mode, and a statistically useful sample of starburst galaxies for this study is not available at present. Future space missions like the Space InfraRed Telescope Facility (SIRFT) and Infrared Imaging Surveyor (IRIS), and the Stratospheric Observatory for Infrared Astronomy (SOFIA) are expected to provide the necessary data base.  The existence of this broad 22 $\\mu$m emission feature complicates the dust model used in the study of the spectral energy distribution of starburst galaxies. Dust grains like graphite, amorphous carbon, silicates, and polycyclic aromatic hydrocarbons may not be representative of all the dust properties in starburst galaxies. Particularly, this broad 22 $\\mu$m emission feature could have significant effects in the derivation of the dust color temperature based on the 20 -- 30 $\\mu$m photometric flux (e.g., the Infrared Astronomical Satellite 25 $\\mu$m data) as well as the number counts of deep surveys in the infrared spectral range to be carried out by SIRTF and IRIS observations, and must be taken into account appropriately.  Arendt et al. (1999) suggested that the carrier of the 22 $\\mu$m feature observed in Cas A is Mg protosilicate based on the good agreement between the observed feature shape and the laboratory spectrum of the Mg protosilicate taken by Dorschner et al. (1980). They found that FeO can also give a good fit to their observed 22 $\\mu$m feature, but the required dust temperature higher than expected and the deficient of emission at wavelegths longer than 30 $\\mu$m led them to rule it out as a promising candidate. If the identification of Mg protosilicate is true, it is the second silicate grain besides the astronomical silicates found in the ISM. More observations are needed to confirm (or test) the suggested identification. Observing the 22 $\\mu$m feature in a variety of astronomical environments will provide useful information in studies on chemical composition and emission mechanism of the carrier.  The major results of this Letter are: (1) a broad 22 $\\mu$m emission dust feature is detected in \\ion{H}{2} regions and starburst galaxies; (2) the 22 $\\mu$m emission feature is similar in shape with the emission feature of newly synthesized dust observed in the ejecta of Cas A, and both of these features arise from the same carrier; and (3) supernovae are probably the major production source of this new interstellar dust."
    },
    "0002/astro-ph0002384_arXiv.txt": {
        "abstract": "Although the population of luminous quasars rises and falls over a period of $\\sim~10^9$ years, the typical lifetime of individual quasars is uncertain by several orders of magnitude.  We show that quasar clustering measurements can substantially narrow the range of possible lifetimes with the assumption that luminous quasars reside in the most massive host halos.  If quasars are long-lived, then they are rare phenomena that are highly biased with respect to the underlying dark matter, while if they are short-lived they reside in more typical halos that are less strongly clustered.  For a given quasar lifetime, we calculate the minimum host halo mass by matching the observed space density of quasars, using the Press-Schechter approximation. We use the results of Mo \\& White to calculate the clustering of these halos, and hence of the quasars they contain, as a function of quasar lifetime. A lifetime of $\\tq = 4 \\times 10^7$ years, the $e$-folding timescale of an Eddington luminosity black hole with accretion efficiency $\\epsilon = 0.1$, corresponds to a quasar correlation length $r_0 \\approx 10 \\hmpc$ in low-density cosmological models at $z = 2 - 3$; this value is consistent with current clustering measurements, but these have large uncertainties. High-precision clustering measurements from the 2dF and Sloan quasar surveys will test our key assumption of a tight correlation between quasar luminosity and host halo mass, and if this assumption holds then they should determine $\\tq$ to a factor of three or better. An accurate determination of the quasar lifetime will show whether supermassive black holes acquire most of their mass during high-luminosity accretion, and it will show whether the black holes in the nuclei of typical nearby galaxies were once the central engines of high-luminosity quasars. ",
        "introduction": "Mounting evidence for the existence of supermassive black holes in the centers of nearby galaxies \\citep[recently reviewed by, e.g.,][]{richstone98} supports the long-standing hypothesis that quasars are powered by black hole accretion \\citep[e.g.,][]{salpeter64,zeldovich64,lyndenbell69}. However, one of the most basic properties of quasars, the typical quasar lifetime $\\tq$, remains uncertain by orders of magnitude. The physics of gravitational accretion and radiation pressure provides one natural timescale, the $e$-folding time $t_e=M_{\\rm BH}/\\dot{M} = 4 \\times 10^8\\, \\epsilon\\, l$ years of a black hole accreting mass with a radiative efficiency $\\epsilon=L/\\dot{M}c^2$ and shining at a fraction $l=L/L_E$ of its Eddington luminosity \\citep{salpeter64}. But while $\\epsilon \\sim 0.1$ and $l \\sim 1$ are plausible values for a quasar, it is possible that black holes accrete much of their mass while radiating at much lower efficiency, or at a small fraction of $L_E$. The task of determining $\\tq$ must therefore be approached empirically.  The observed evolution of the quasar luminosity function imposes a strong upper limit on $\\tq$ of about $10^9$ years, since the whole quasar population rises and falls over roughly this interval \\citep[see, e.g.,][]{osmer98}. The lifetime of individual quasars could be much shorter than the lifetime of the quasar population, however, and lower limits of $\\tq \\sim 10^5$ years rest on indirect arguments, such as the requirement that quasars maintain their ionizing luminosity long enough to explain the proximity effect in the Ly$\\alpha$ forest \\citep[e.g.,][]{bajtlik88,bechtold94}. A typical lifetime $\\tq \\sim 10^9$ years would imply that quasars are rare phenomena, arising in at most a small fraction of high-redshift galaxies.  Conversely, a lifetime as low as $\\tq \\sim 10^5$ years would imply that quasars are quite common, suggesting that a large fraction of present-day galaxies went through a brief quasar phase in their youth.  The comoving space density $\\Phi(z)$ of active quasars at redshift $z$ is proportional to $\\tq n_H(z)$, where $n_H$ is the comoving space density of quasar hosts.  ``Demographic'' studies of the local black hole population \\citep[e.g.,][]{mag98,salucci99,marel99} have opened up one route to determining the typical quasar lifetime: counting the present-day descendants of the quasar central engines in order to estimate $n_H(z)$ and thus constrain $\\tq$ by matching $\\Phi(z)$. Roughly speaking, the ubiquity of black holes in nearby galaxies suggests that quasars are common and that $\\tq$ is likely in the range $10^6$ -- $10^7$ \\citep[e.g.,][]{richstone98,haehnelt98,salucci99}. However, as \\citet{richstone98} emphasize, the lifetime estimated in this way depends crucially on the way one links the mass of a present-day black hole to the luminosity of a high-redshift quasar, which in turn depends on assumptions about the growth of black hole masses since the quasar epoch via mergers or low-efficiency accretion.  In this paper we propose an alternative route to the quasar lifetime, using measurements of high-redshift quasar clustering. The underlying idea goes back to the work of \\citet{kaiser84} and \\citet{bbks86}: in models of structure formation based on gravitational instability of Gaussian primordial fluctuations, the rare, massive objects are highly biased tracers of the underlying mass distribution, while more common objects are less strongly biased. Therefore, a longer quasar lifetime $\\tq$ should imply a more clustered quasar population, provided that luminous quasars reside in massive hosts. The specific calculations that we present in this paper use the Press-Schechter (1974; hereafter PS) approximation for the mass function of dark matter halos and the \\citet[hereafter MW]{mw96} and \\citet{jing98} approximations for the bias of these halos as a function of mass.  The path from clustering to quasar lifetime has its own uncertainties; in particular, our predictions for quasar clustering will rely on the assumption that the luminosity of a quasar during its active phase is a monotonically increasing function of the mass of its host dark matter halo.  However, the assumptions in the clustering approach are at least very different from those in the black hole mass function approach, and they can be tested empirically by detailed studies of quasar clustering as a function of luminosity and redshift.  Our theoretical model of quasar clustering follows a general trend in which the study of quasar activity is embedded in the broader context of galaxy formation and gravitational growth of structure \\citep[e.g.,][]{efstathiou88,turner91,haehnelt93,katz94,haehnelt98,haiman98, monaco00,kauffmann00}. This paper also continues a theme that is prominent in recent work on the clustering of Lyman-break galaxies, namely that the clustering of high-redshift objects is a good tool for understanding the physics of their formation and evolution \\citep[e.g.,][]{adelberger98,katz99,kolatt99,mmw99}. Our model of the quasar population is idealized, but by focusing on a simple calculation with clearly defined predictions, we hope to highlight the link between quasar lifetime and clustering strength. After presenting the theoretical results, we will draw some inferences from existing estimates of the quasar correlation length. However, our study is motivated mainly by the anticipation of vastly improved measurements of quasar clustering from the 2dF and Sloan quasar surveys \\citep[see, e.g.,][]{boyle99,fan99,york00}. These measurements can test various hypotheses about the origin of quasar activity, including our primary assumption of a monotonic relation between quasar luminosity and host halo mass.  If this assumption proves valid, then the first major physical result to emerge from the 2dF and Sloan measurements of high-redshift quasar clustering will be a new determination of the typical quasar lifetime. ",
        "conclusions": "\\label{sec:dis}  \\subsection{Sensitivity to model details} \\label{sec:sen}  As already mentioned in \\S\\ref{sec:life}, the definition of a ``halo lifetime'' is somewhat ambiguous.  We have so far adopted a definition of $t_H$ as the median time before a halo of mass $M$ is incorporated into a halo of mass $2M$.  If we increase this mass ratio from 2 to 5 (a rather extreme value), then the typical halo lifetimes in our CDM models increase by factors of $2-4$. Since it is the ratio $t_Q/t_H$ that enters our determination of $\\mmin$ (eq.~[\\ref{eqn:phimatch2}]), and hence fixes the bias factor, this change in $t_H$ would require an equal increase in $t_Q$ to maintain the same clustering length $r_1$. We conclude that the ambiguity in halo lifetime definition introduces a factor $\\sim 2$ uncertainty in the determination of $t_Q$ from clustering measurements, in the context of our model.  We have also assumed that quasar luminosity is perfectly correlated with host halo mass, so that matching the space density of an absolute-magnitude limited sample imposes a sharp cutoff in the host mass distribution at $M=\\mmin$.  If there is some scatter in the luminosity--host mass relation, then some halos with $M<\\mmin$ will host a quasar above the absolute-magnitude limit and some halos with $M>\\mmin$ will not.  We can model such an effect by introducing a soft cutoff into equation~(\\ref{eqn:phimatch2}): \\begin{equation} \\Phi(z) = \\int_{0}^{\\infty} dM\\;g(M)\\;\\frac{\\tq}{t_H(M, z)} n(M, z) \\end{equation} with \\begin{equation} g(M) = \t\\left\\{ \\begin{array}{ll} 0       & {\\rm for} \\; M < \\frac{\\mmin}{\\alpha} \\\\ \\left( \\frac{\\alpha}{\\mmin (\\alpha^2 - 1)} \\right) M - \\frac{1}{\\alpha^2 - 1}    & {\\rm for} \\; \\frac{\\mmin}{\\alpha} < M < \\alpha \\mmin \\\\ 1       & {\\rm for} \\; M > \\alpha \\mmin \\end{array} \\right. \\label{eqn:cut} \\end{equation} and $\\alpha > 1$.  Adopting a soft cutoff slightly decreases $\\mmin$ and, more significantly, reduces the value of $\\beff$ by allowing some quasars to reside in lower mass halos, which are less strongly biased. Quantitatively, we find that setting $\\alpha = 2$, which corresponds to including halos down to $M=\\mmin/2$, decreases the clustering length by $\\la 6$\\% for the shortest quasar lifetimes and $\\la 10$\\% for the longest quasar lifetimes.  Matching a fixed $r_1$ with an $\\alpha=2$ cutoff requires lifetimes that are longer by a factor $\\sim 1-1.5$ at short $\\tq$ and $\\sim 2-2.5$ at long $\\tq$. Longer lifetimes are more sensitive to scatter in the luminosity-host mass relation because $\\beff$ depends more strongly on $\\mmin/M_*$ for these rarer objects. The assumption of a perfectly monotonic relation between quasar luminosity and host mass leads to the smallest $\\tq$ for a given $r_1$. Thus if any scatter does exist in this relation, our model predictions for $\\tq$ effectively become lower limits to the quasar lifetime.  Another simplification of our model is the assumption that a quasar is either ``on'' or ``off'' -- each quasar shines at luminosity $L$ for time $t_Q$, perhaps divided among several episodes of activity, and the rest of the time it is too faint to appear in a luminous quasar sample. More realistically, variations in the accretion rate and radiative efficiency will cause the quasar luminosity to vary, especially if the black hole mass itself grows significantly during the luminous phase. Nonetheless, the maximum luminosity will still depend on the maximum black hole mass. At a given time, the luminous quasar population will include black holes shining at close to their maximum luminosity and ``faded'' black holes of higher mass. Because the host halos lie on the steeply falling tail of the mass function, the first component of the population always dominates over the second, and we therefore expect our clustering method to yield the time $t_Q$ for which a quasar shines within a factor $\\sim 2$ of its peak luminosity. More strongly faded quasars are too rare to make much difference to the space density or effective bias.  To illustrate this point, we consider the model of \\citet{haehnelt98} in which a quasar hosted by a halo of mass $M$ has a luminosity history $L(t) = L_0(M)\\,{\\rm exp}(-t/t_Q)$, with a maximum luminosity $L_0(M) = \\alpha\\,M$ proportional to the halo mass. In this model, the time that a quasar shines above the luminosity threshold $L_{\\rm min} = L_0 (\\mmin)$ of a survey is the visibility time $t_Q' = t_Q\\,{\\rm ln}(M/\\mmin)$. We can calculate $\\mmin$ for a given space density by substituting $t_Q'$ for $t_Q$ in equation~(\\ref{eqn:phimatch2}), then calculate $\\beff(\\mmin)$ by multiplying the integrands in the numerator and denominator of equation~(\\ref{eqn:beff}) by the visibility weighting factor ${\\rm ln}(M/\\mmin)$. The middle curves in Figure~\\ref{fig:lt} compare $r_1(t_Q)$ for the on--off ({\\it solid line}) and exponential decay ({\\it dotted line}) models, in the case of $\\Lambda$CDM at $z = 3$ with our standard $\\Phi(z)$. The curves are remarkably similar, showing that the lifetime inferred from clustering assuming an on--off model would be close to the $e$-folding timescale in an exponential decay model. The curves for the exponential decay model are slightly shallower because at low $\\mmin$ (low $t_Q$) the mass function is not as steep, allowing faded quasars in more massive halos to make a larger contribution to $\\beff$ and thereby raise $r_1$. Although results for a different functional form of $L(M,t)$ would differ in detail, we would expect the lifetime inferred from clustering to be close to the ``half-maximum'' width of the typical luminosity history, for the general reasons discussed above.  \\begin{figure*}[t] \\centerline{ \\epsfxsize=3.5truein \\epsfbox[65 165 550 730]{martini.f8.eps} } \\caption{ \\footnotesize Clustering length vs. $\\tq$ for the $\\Lambda$CDM model at $z=3$ for two different models of quasar luminosity evolution at three different space densities $\\Phi(z)$. The ``on -- off'' model ({\\it solid lines}) assumes the quasar luminosity is constant throughout its lifetime $\\tq$ and is the standard model we discuss in this paper. The central line shows results for the $\\Lambda$CDM model at $z = 3$ with our standard $\\Phi(z)$. The other solid curves, related to the first by simple horizontal shifts, show results for space densities different by factors of $10$ and $1/10$ (bottom and top). In the exponential model ({\\it dotted lines}), the quasar luminosity starts at some maximum luminosity proportional to the halo mass and decays with an $e$-folding timescale $\\tq$. The middle line again corresponds to our standard $\\Phi(z)$, and the other two dotted curves show results for space densities different by factors of $10$ and $1/10$ (bottom and top). } \\label{fig:lt} \\end{figure*}  As mentioned in \\S\\ref{sec:over}, we assume that quasars radiate isotropically. If they radiate instead with an average beaming factor $f_B < 1$, then the true value of $\\Phi(z)$ is larger than the observed value by a factor $f_B^{-1}$. The implied lifetime for a given $r_1$ would therefore be larger by a factor $f_B^{-1}$ as well.   \\subsection{Interpretation of Existing Data} \\label{sec:data}  After several attempts \\citep{osmer81, webster82}, quasar clustering was first detected by \\citet{shaver84}, and later by \\citet{shanks87} and \\citet{iovino88}.  However, measurements of quasar clustering are still hampered by small, sparse samples, and even the best studies to date yield detections with only several-$\\sigma$ significance. Given the limitations of current data, it is not surprising that different authors reach different conclusions about the strength of clustering and its evolution.  Analyzing a combined sample of quasars with $0.3<z<2.2$ from the Durham/AAT UVX Survey, the CFHT survey, and the  Large Bright Quasar Sample, \\citet{shanks94} and \\citet{croom96} find a reasonable fit to the data with an $\\om=1$ model that has $\\xi(r)=(r/r_0)^{-\\gamma}$, $\\gamma \\approx 1.8$, and a constant comoving correlation length $r_0=6\\hmpc$. \\citet{lafranca98} report a higher correlation length, $r_0=9.1 \\pm 2.0\\hmpc$ for a $\\gamma=1.8$ power law, in their $1.4 < z < 2.2$ sample.  If we adopt $r_0 \\approx 8\\hmpc$ at $z=2$ and a corresponding $r_1 \\approx 11\\hmpc$, then the implied quasar lifetime is $\\sim 10^{7.5}$ years for the $\\tau$CDM model and $\\sim 10^8$ years for SCDM.  The $r_0$ values quoted above are for $\\om=1$, and because quasar pair separations are measured in angle and redshift, they should be increased by a factor $\\sim 1.5$ in an $\\om=0.3$, $\\ol=0.7$ universe and a factor $\\sim 1.3$ in an $\\om=0.3$, $\\ol=0$ universe (roughly the inverse cube-roots of the volume ratios listed in Table A1).  Adopting $r_1 \\approx 16\\hmpc$ implies a lifetime $t_Q \\sim 10^{7-7.5}$ years in our low-density models. However, these numbers must be considered highly uncertain because of the limitations of current data and because the space densities of the various observational samples do not necessarily match those assumed in our model predictions.  All of these measurements are based mainly on quasars with $z<2$.  At higher redshift, \\citet{kundic97} and \\citet{stephens97} have investigated clustering in the Palomar Transit Grism Survey (PTGS; \\citealt{schneider94}). Fitting a $\\gamma=1.8$ power law, \\citet{stephens97} find $r_0=17.5 \\pm 7.5\\hmpc$ for $z>2.7$.  This high correlation length (inferred from the presence of three close pairs in a sample of 90 quasars) could be a statistical fluke, but in the context of our model it is tempting to see it as a consequence of the high luminosity threshold of the PTGS survey, which might lead it to pick out the most strongly clustered members of the quasar population.  \\subsection{Prospects} \\label{sec:pros}  The 2dF \\citep{boyle00,shanks00} and Sloan \\citep{york00} quasar surveys will transform the study of quasar clustering over the next several years, yielding high-precision measurements for a wide range of redshifts.  These measurements will allow good determination of the typical quasar lifetime $\\tq$ in the context of the model presented here.  They will also test the key assumption of this model, the monotonic relation between quasar luminosity and host halo mass, by characterizing the clustering as a function of redshift and, especially, as a function of quasar absolute magnitude.  Figure~\\ref{fig:lt} illustrates this test for the $\\Lambda$CDM model at $z=3$. Brighter quasars have a lower space density $\\Phi(z)$, so they should have a higher minimum host halo mass $\\mmin$, and, because of the higher bias of more massive halos, they should exhibit stronger clustering.  Fainter, more numerous quasars should exhibit weaker clustering.  Figure~\\ref{fig:lt} shows the predicted $r_1$ vs. $\\tq$ relation for samples with 1/10 and 10 times the space density of our standard case (3.42 quasars per square degree per unit redshift; see Table~\\ref{tbl:qlf}). In our standard on--off model ({\\it solid lines}), a change in $\\Phi(z)$ in equation~(\\ref{eqn:phimatch2}) can be exactly compensated by changing $\\tq$ by the same factor, so the solid curves in Figure~\\ref{fig:lt} are simply shifted horizontally relative to each other.  Our predictions in Figure~\\ref{fig:r1cdm} (see eq.~[\\ref{eqn:fits}]) can therefore be transformed to any quasar space density by changing $\\tq$ in proportion to $\\Phi(z)$. In the exponential decay model ({\\it dotted lines}), the scaling of $t_Q$ with $\\Phi(z)$ is no longer exact, though it is still a good approximation.  If there is a large dispersion in the relation between quasar luminosity and host halo mass, then the dependence of clustering strength on quasar space density will be much weaker than Figure~\\ref{fig:lt} predicts. Detection of the predicted trend between luminosity and clustering, or definitive demonstration of its absence, would itself provide an important insight into the nature of quasar host halos. More generally, the parameters of a model that incorporates scatter (such as the $\\alpha$ prescription of equation~[\\ref{eqn:cut}]) could be determined by matching the observed relation between $r_1$ and $\\Phi(z)$.  If the observations do support a tight correlation between luminosity and host halo mass, then the first property of quasars to emerge from the 2dF and Sloan clustering studies will be the typical lifetime $\\tq$. For the low-density models in Figure~\\ref{fig:r1cdm}, the slope of the correlation between $r_1$ and $\\log_{10}\\tq$ is $\\sim 10$, so a determination of $r_1$ with a precision of $2\\hmpc$ would constrain $\\tq$ to a factor of $10^{0.2} \\approx 1.6$, for a specified cosmology.  By the time these quasar surveys are complete, a variety of observations may have constrained cosmological parameters to the point that they contribute negligible uncertainty to this constraint.  Instead, the uncertainty in $\\tq$ will probably be dominated by the limitations of the quasar population model, e.g., the approximate nature of the assumptions that the quasar luminosity tracks the halo mass, that there is only quasar per halo, and that the average lifetime $\\tq$ is independent of quasar luminosity. These assumptions can be tested empirically to some degree, but not perfectly. Despite these limitations, it seems realistic to hope that $\\tq$ can be constrained to a factor three or better by high-precision clustering measurements, a vast improvement over the current situation. It is worth reiterating that our assumption of a perfectly monotonic relation between luminosity and halo mass leads to the smallest $\\tq$ for an observed $r_1$, since with a shorter lifetime there are simply not enough massive, highly biased halos to host the quasar population.  A determination of $\\tq$ to a factor of three will be sufficient to address fundamental issues about the physics of quasars and galactic nuclei.  Comparison of $\\tq$ to the Salpeter timescale will answer one of the most basic questions about supermassive black holes: do they shine as they grow? If $\\tq \\ga 4\\times 10^7$ years, the $e$-folding timescale for $L \\sim L_E$, $\\epsilon \\sim 0.1$, then quasar black holes increase their mass by a substantial factor during their optically bright phase. If $\\tq$ is much shorter than this, then the black holes must accrete most of their mass at low efficiency, or while shining at $L \\ll L_E$. A short lifetime could indicate an important role for advection dominated accretion (\\citet{narayan98} and references therein), or it could indicate that black holes acquire much of their mass through mergers with other black holes, emitting binding energy in the form of gravitational waves rather than electromagnetic waves. A determination of $\\tq$ would also resolve the question of whether the black holes in the nuclei of local galaxies are the remnants of dead quasars. For example, \\citet{richstone98} infer a lifetime $t_Q \\sim 10^6$ years by matching the space density of local spheroids that host black holes of mass $M \\ga 4\\times 10^8 M_\\odot$ to the space density of high-redshift quasars of luminosity $L_E(M) \\ga 6 \\times 10^{46}\\;{\\rm erg}\\;{\\rm s}^{-1}$. If clustering implies a much longer lifetime, then these numerous local black holes may once have powered active nuclei, but they were not the engines of the luminous, rare quasars.  We have assumed in our model that quasar activity is a random event in the life of the parent halo. Quasar activity might instead be triggered by a major merger, by a weaker ``fly-by'' interaction, or by the first burst of star formation in the host galaxy. Regardless of the trigger mechanism, the lifetime will be the dominant factor in determining the strength of high-redshift quasar clustering, if our assumed link between luminosity and halo mass holds.  However, different triggering mechanisms might be diagnosed by more subtle clustering properties, such as features in the correlation function at small separations, or higher-order correlations.  At low redshift, where the evolution of the quasar population is driven by fueling rather than by black hole growth, the nature of the triggering mechanism might play a major role in determining quasars' clustering properties. The calculations presented here illustrate the promise of quasar clustering as a tool for testing ideas about quasar physics, a promise that should be fulfilled by the large quasar surveys now underway."
    },
    "0002/astro-ph0002498_arXiv.txt": {
        "abstract": "The evolution of the star formation rate in the Galaxy is one of the key ingredients quantifying the formation and determining the chemical and luminosity evolution of galaxies. Many complementary methods exist to infer the star formation history of the components of the Galaxy, from indirect methods for analysis of low-precision data, to new exact analytic methods for analysis of sufficiently high quality data. We summarise available general constraints on star formation histories, showing that derived star formation rates are in general comparable to those seen today.  We then show how colour-magnitude diagrams of volume- and absolute magnitude-limited samples of the solar neighbourhood observed by Hipparcos may be analysed, using variational calculus techniques, to reconstruct the local star formation history. The remarkable accuracy of the data coupled to our maximum-likelihood variational method allows objective quantification of the local star formation history with a time resolution of $\\approx 50$Myr.  Over the past 3Gyr, the solar neighbourhood star formation rate has varied by a factor of $\\sim$4, with characteristic timescale about 0.5Gyr, possibly triggered by interactions with spiral arms. ",
        "introduction": "Star formation rates and histories can be estimated in special cases from a combination of chemical evolution models and the total stellar mass formed into stars. Basically, this exploits the stellar evolutionary, and Type~I supernova timescales for element production. One requires that stars formed at a rate consistent with their chemical distributions, whether a $\\delta$-function, a range in the products of both Type~I and Type~II supernaova, products of single supernovae with time for efficient mixing of the ISM, or whatever. Combining these albeit crude estimates of the {\\sl duration} of star formation with a calculation of the stellar mass formed, provides a star formation rate. Perhaps surprisingly, given the crude calculation, such derived rates are both similar to those determined more accurately today, and are all quite low.    \\begin{table}[h] \\begin{center} \\caption{A summary of star formation rates, and durations of star formation, in some Galactic stellar populations. These values are derived from combination of chemical element scatter and masses.} \\vskip 10pt \\leavevmode \\footnotesize \\begin{tabular}[tbh]{|l|c|c|} \\hline && \\\\[-4pt] { Population }  & Duration & Formation Rate \\\\ \\hline && \\\\[-4pt] &  (years) & ${\\mathcal M}_{\\odot}yr^{-1}$ \\\\ \\hline && \\\\[-4pt] Globular cluster & $\\le10^8$ & $\\ge 0.01$ \\\\ $\\omega$Cen&$\\ge10^8$ & $\\le 0.1$ \\\\ Halo, [Fe/H]$\\le-2.0$ & $\\le10^8$ & $\\sim 1$ \\\\ Halo, [Fe/H]$\\sim-1.5$ & $\\le10^9$ &  $\\sim 1$ \\\\ Bulge; high [$\\alpha$/Fe] & few.10$^8$ & 10-100 \\\\ Bulge; low [$\\alpha$/Fe] & few.10$^9$ & 10-100 \\\\ Thick Disk & few.10$^9$ & 1-10 \\\\ Current Disk & $10^{10}$ & $\\sim 1-10$ \\\\ Inner Disk & ? & ? \\\\ Satellite dSph & many.10$^9$ & $\\le 10^{-3}$ \\\\ \\hline && \\\\[-4pt] Assembly & early & \\\\ \\hline && \\\\[-4pt] Infall & continuing? & $\\sim 4$Gyr \\\\ \\hline \\end{tabular} \\vspace{-0.5cm} \\end{center} \\end{table}  In cases where no spread in element ratios is seen, and there is no range in [Fe/H], star formation was plausibly complete before new chemical elements could be produced; perhaps globular clsuters are an example of this case. For field halo stars with [Fe/H]$\\ge -2$, where a wide range in [Fe/H] but a very small range in [$\\alpha$/H] is seen, star formation must have continued for long enough for efficient mixing of supernova ejecta into the ISM. Since no products of Type~I supernovae are seen, this brackets allowed star formation durations. By applying such qualitative considerations, we deduce that most of the Milky way formed at a star formation rate which is comparable to that of today. Only the Galactic Bulge, and the inner disk, where star formation histories remain very poorly known, are available to retrieve the Milky Way's place as a `typical object' on the Madau plot. These simple calculations are summarised in Table~(1). ",
        "conclusions": "Star formation rates and histories can be estimated in special cases from a combination of chemical evolution models and the total stellar mass formed into stars. Basically, this exploits the stellar evolutionary, and Type~I supernova timescales for element production. In cases where no spread in element ratios is seen, and there is no range in [Fe/H], star formation was plausibly complete before new chemical elements could be produced; perhaps globular clsuters are an example of this case. For field halo stars with [Fe/H]$\\ge -2$, where a wide range in [Fe/H] but a very small range in [$\\alpha$/H] is seen, star formation must have continued for long enough for efficient mixing of supernova ejects into the ISM. Since no products of Type~I supernovae are seen, this brackets allowed star formation durations. By applying such qualitative considerations, we deduce that most of the Milky way formed at a star formation rate which is comparable to that of today. Only the Galactic Bulge, and the inner disk, where star formation histories remain very poorly known, are available to retrieve the Milky Way's place as a `typical object' on the Madau plot.  In the immediate Solar neighbourhood, and in dSph satellite galaxies, the high quality data and the simple stellar populations respectively allow us to be more quantitative.  We have applied the objective variational calculus method for the reconstruction of star formation histories from observed colour magnitude data, developed in our Paper~(I), to the data in the Hipparcos catalogue, yielding the star formation history of the solar neighbourhood over the last 3 Gyr. Surprisingly, a structured star formation history is obtained, showing a cyclic pattern with a period of about 0.5 Gyr, superimposed on some underlying star formation activity which increases slightly with age.  No random bursting behaviour was found at the time resolution of 0.05 Gyr of our method. A first order density wave model for the repeated encounter of galactic arms could explain the observed regularity."
    },
    "0002/astro-ph0002451_arXiv.txt": {
        "abstract": "We show that the Tully-Fisher relation observed for spiral galaxies can be explained in the current scenario of galaxy formation without invoking subtle assumptions, provided that galactic-sized dark haloes have low concentrations which do not change significantly with halo circular velocity. This conclusion does not depend significantly on whether haloes have cuspy or flat profiles in the inner region. In such a system, both the disk and the halo may contribute significantly to the maximum rotation of the disk, and the gravitational interaction between the disk and halo components leads to a tight relation between the disk mass and maximum rotation velocity. The model can therefore be tested by studying the Tully-Fisher zero points for galaxies with different disk mass-to-light ratios. With model parameters (such as the ratio between disk and halo mass, the specific angular momentum of disk material, disk formation time) chosen in plausible ranges, the model can well accommodate the zero-point, slope, and scatter of the observed Tully-Fisher relation, as well as the observed large range of disk surface densities and sizes. In particular, the model predicts that low surface-brightness disk galaxies obey a Tully-Fisher relation very similar to that of normal disks, if the disk mass-to-light ratio is properly taken into account. About half of the gravitational force at maximum rotation comes from the disk component for normal disks, while the disk contribution is lower for galaxies with a lower surface density.  The halo profile required by the Tully-Fisher relation is as concentrated as that required by the observed rotation curves of faint disks, but less concentrated than that given by current simulations of CDM models. We discuss the implication of such profiles for structure formation in the universe and for the properties of dark matter. Our results cannot be explained by some of the recent proposals for resolving the conflict between conventional CDM models and the observed rotation-curve shapes of faint galaxies. If dark matter self-interaction (either scattering or annihilation) is responsible for the shallow profile, the observed Tully-Fisher relation requires the interaction cross section $\\sigma_X$ to satisfy $\\langle\\sigma_{X}\\vert v\\vert\\rangle/m_{X} \\sim 10^{-16} {\\rm cm^{3}\\,s^{-1}\\,GeV^{-1}}$, where $m_X$ is the mass of a dark matter particle. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002445_arXiv.txt": {
        "abstract": "We report results of a high-resolution imaging search for the galaxy associated with the damped Lyman-$\\alpha$ (DLA) absorber at $z=1.892$ toward the $z_{em}=2.543$ quasar LBQS 1210+1731, using HST/NICMOS. The images were obtained in the broad filter F160W and the narrow filter F190N with camera 2 on NICMOS, and were aimed at detecting the absorber in the rest-frame optical continuum and in H-$\\alpha$ line emission from the DLA absorber. After suitable point spread function (PSF) subtractions, a feature is seen in both the broad-band and narrow-band images, at a projected separation of 0.25$\\arcsec$ from the quasar. This feature may be associated with the DLA absorber, although we cannot completely rule out that it could be a PSF artifact. If associated with the DLA, the object would be $\\approx 2-3$ $h_{70}^{-1}$ kpc in size with a flux of $9.8 \\pm 2.4$ $\\mu$Jy in the F160W filter, implying a luminosity at $\\lambda_{central}=5500$ {\\AA} in the rest frame of $1.5 \\times 10^{10}$ $h_{70}^{-2}$ L$_{\\odot}$ at $z=1.89$, for $q_{0}=0.5$. However, a comparison of the fluxes in the broad and narrow filters indicates that most of the flux in the narrow-band filter is continuum emission, rather than red-shifted H-$\\alpha$ line emission. This suggests that if this object is the absorber, then either it has a low star formation rate (SFR), with a 3 $\\sigma$ upper limit of 4.0 $h_{70}^{-2}$ M$_{\\odot}$ yr$^{-1}$, or dust obscuration is important. It is possible that the H-$\\alpha$ emission may be extinguished by dust, but this seems unlikely, given the typically low dust-to-gas ratios observed in DLAs. Alternatively, the object, if real, may be associated with the host galaxy of the quasar rather than with the damped Ly-$\\alpha$ absorber. H-band images obtained with the NICMOS camera 2 coronagraph show a much fainter structure $\\approx 4-5$ $h_{70}^{-1}$ kpc in size and containing four knots of continuum emission, located 0.7$\\arcsec$ away from the quasar. This structure is not seen in images of comparison stars after similar PSF subtractions, and is also likely to be associated with the absorbing galaxy or its companions, although we do not know its redshift.  We have probed regions far closer to the quasar sight-line than in most previous studies of high-redshift intervening DLAs. The two objects we report mark the closest detected high-redshift DLA candidates yet to any quasar sight line. If the features in our images are associated with the DLA, they suggest faint, compact, somewhat clumpy objects rather than large, well-formed proto-galactic disks or spheroids. If the features are PSF artifacts, then the constraints on sizes and star-formation rates of the DLA are even more severe. The size, luminosity, and SFR estimates mentioned above should therefore be conservatively considered as upper limits. ",
        "introduction": "Damped Lyman-$\\alpha$ absorption systems detected in spectra of high-redshift quasars are believed to be the progenitors of present-day galaxies, because they show high H~I column densities (log $N_{HI} \\ge 20.0$) and display absorption lines of several heavy elements. However, there are various competing ideas regarding the nature of the galaxies underlying the DLAs. Wolfe et al. (1986) suggested that the DLAs are rotating proto-disks. This suggestion has also been made by Prochaska \\& Wolfe (1997, 1998), based on asymmetric line profiles of the heavy-element absorption lines in DLAs. On the other hand, gas-rich dwarf galaxies have also been suggested as candidate objects for the DLAs (York et al. 1986; Matteucci, Molaro, \\& Vladilo 1997). Recently, Jimenez, Bowen, \\& Matteucci (1999) have suggested that high-redshift DLAs may arise in low-surface brightness galaxies. The lack of substantial chemical evolution found in studies of element abundances in DLAs (e.g., Pettini et al. 1999; Kulkarni, Bechtold, \\& Ge 2000a) also shows that the currently known population of DLAs seems to be dominated by metal-poor objects, so DLAs may consist of dwarf or low-surface brightness galaxies with modest star formation rates. Unfortunately, it is hard to determine what type of galaxies underlie the DLAs, since most previous efforts to directly image the high-redshift DLAs have failed. A few detections have been made at low redshifts, which showed those DLAs to arise in low surface-brightness galaxies (see, e.g., Steidel et al. 1995a, 1995b; LeBrun et al. 1997). But high-redshift DLAs with $z_{abs} < z_{em}$ have proven hard to detect, and the question of the nature of galaxies giving rise to these DLAs is still open.  Many of the previous attempts to detect the emission from DLAs concentrated on the Ly-$\\alpha$ emission, which is an expected signature from a star-forming region (e.g. Smith et al. 1989; Hunstead, Pettini, \\& Fletcher 1990; Lowenthal et al. 1995). There have been only a few Ly-$\\alpha$ detections of DLAs so far. M{\\o}ller \\& Warren (1998) and M{\\o}ller, Warren, \\& Fynbo (1998) detected Ly-$\\alpha$ emission in the fields of two DLAs at $z=2.81$ and $z=1.93$. However, both of these DLAs have $z_{abs} \\approx z_{em}$ and may be different from the general population of intervening DLAs. Djorgovski et al. (1996) and Djorgovski (1997) reported Ly-$\\alpha$ emitting objects with $R \\sim 25$ (and inferred SFRs of a few M$_{\\odot}$ yr$^{-1}$) in fields of a few DLAs, located at 2-3 $\\arcsec$ from the quasar. However, the Ly-$\\alpha$ technique cannot definitively measure the star formation rates of the DLAs because of the generally unquantifiable effects of dust extinction in the systems. The lack of detections in the other Ly-$\\alpha$ studies of interevening DLAs could indicate either that DLAs have low star formation rates (SFR) or that the emission is extinguished by dust. As pointed out by Charlot \\& Fall (1991), even small quantities of dust are sufficient to extinguish the Ly-$\\alpha$ emission, since resonant scattering greatly increases the path length of Ly-$\\alpha$ photons attempting to escape from an H I cloud. Indeed, observations of reddening of background quasars and evidence for depletion of Cr, Fe, Ni etc. relative to Zn suggest the presence of a small amount of dust in DLAs (see, e.g., Pei, Fall, \\& Bechtold 1991; Pettini et al. 1997; Kulkarni, Fall, \\& Truran 1997). Thus, it is hard to constrain the SFRs in DLA galaxies using the non-detections or weak detections of Ly-$\\alpha$ emission.  The issues of dust and SFR in high redshift DLAs are also important in view of recent claims based on mid-IR and far-IR observations that a large fraction of the star formation at high redshifts is hidden from us by dust obscuration (e.g., Elbaz et al. 1998; Clements et al. 1999). One way to discern whether the previous non-detections of Ly-$\\alpha$ were due to low SFR or presence of dust is to search for longer wavelength emission lines less affected by dust extinction and not subject to resonant scattering.  The ground-based near-IR spectroscopic survey of Bunker et al. (1999), which searched for redshifted H-$\\alpha$ emission in $11 \\arcsec \\times 2.5 \\arcsec$ regions around 6 quasars with DLAs at $z > 2$ and reached 3 $\\sigma$ detection levels of 6-18 M$_{\\odot}$ yr$^{-1}$, failed to detect any redshifted H-$\\alpha$ emission from the DLAs in their sample. Some of the ground-based narrow-band photometric surveys for H-$\\alpha$ emission from DLAs have also failed to detect any emission line objects in the DLA fields (e.g., Teplitz, Malkan, \\& McLean 1998, who however found H-$\\alpha$ emitters in the fields of some weaker non-DLA metal line systems). Some other narrow-band searches for H-$\\alpha$ emission have revealed multiple objects in the DLA fields separated by several arcseconds from the quasar (2-12$\\arcsec$ for Bechtold et al. 1998, 9-120 $\\arcsec$ for Mannucci et al. 1998). These surveys, which had 3 $\\sigma$ detection limits of $\\sim 5$  M$_{\\odot}$ yr$^{-1}$ (Bechtold et al. 1998) or $\\gtrsim 10$ M$_{\\odot}$ yr$^{-1}$ (Mannucci et al. 1998), found the H-$\\alpha$ emitting objects to have a wide range of inferred SFRs (10-20 M$_{\\odot}$ yr$^{-1}$ for Bechtold et al. 1998, 6-90 M$_{\\odot}$ yr$^{-1}$ for Mannucci et al. 1998). The relatively large separations of these emission line objects from the quasars indicates that they are not the DLA absorbers themselves, but star-forming companions in the same larger structure (e.g. sheet or filament) as the DLA. None of these ground-based surveys has been able to probe the regions very close to the quasar sightline (angular separations $< 2$ $\\arcsec$), because of the limitations imposed by seeing in these studies. While these studies offer interesting information about the environments of the DLAs, high sensitivity diffraction-limited imaging is necessary for the detection of the DLA absorbers themselves (to probe small angular separations), and thus for determining the morphology and SFRs of the DLAs. The HST WFPC2 study of Le Brun et al. (1997) has detected candidates with angular separations $< 2$ $\\arcsec$ in broad band images for six DLAs at $z < 1$ and one DLA at $z=1.78$. However, the information obtained from this study about the nature of high-redshift DLAs is limited since no narrow-band images were obtained and since the sample contained only 1 DLA at $z > 1$. As mentioned earlier, the HST WFPC2 study of M{\\o}ller \\& Warren (1998 and references therein) detected Ly-$\\alpha$ emission in a $z_{abs} > z_{em}$ DLA, but this DLA may differ from intervening DLAs.  To summarize, many previous attempts to detect emission from high-redshift intervening DLAs have failed. The few detections so far consist mainly of either weak Ly-$\\alpha$ detections (which cannot constrain the SFR completely) or detections of H-$\\alpha$ companions at fairly large angular separations from the quasars. There are only four objects detected so far in fields of high-$z$ interevening DLAs at small angular separations. These objects have impact parameters between 4.3 and 11.5 h$_{70}^{-1}$ kpc (where $H_{0} = 70$ h$_{70}$ km s$^{-1}$ Mpc$^{-1}$), and are promising candidates for the DLAs in those sightlines (see M{\\o}ller \\& Warren 1998 and references therein; the other DLA impact parameter data listed in M{\\o}ller \\& Warren 1998 are biased toward $z_{abs} \\approx z_{em}$ DLAs.). To further increase the number of promising candidates for high-redshift intervening DLAs, it is necessary to carry out more deep high spatial resolution near-infrared searches for DLAs.  We have obtained deep diffraction-limited images of three DLAs at $z \\sim 2$ with the Near Infrared Camera and Multi-Object Spectrometer (NICMOS) onboard the Hubble Space Telescope (HST). Here we describe our NICMOS observations of the quasar LBQS 1210+1731 ($z_{em}=2.543 \\pm 0.005$; Hewett, Foltz, \\& Chaffee 1995), which has a spectroscopically known damped Ly-$\\alpha$ absorber ($z_{abs}=1.8920$ and log $N_{HI} = 20.6$; Wolfe et al. 1995). Our observations have the unique benefit of combining high near-IR sensitivity and high spatial resolution with a more stable and quantifiable PSF than is currently possible with ground-based observations. A further feature of some of our observations is the use of the NICMOS coronagraph, which greatly decreases the scattered light background outside of the coronagraphic hole and therefore allows a study of the environment of the DLA. Our analysis indicates two objects at 0.25 $\\arcsec$ and 0.7 $\\arcsec$ from the quasar that we cannot explain as any known artifacts of the PSF. We believe that these objects are likely to be real and may be associated with the DLA and its companions, at impact parameters of 1.5 and 3.8 h$_{70}^{-1}$ kpc. We have thus probed regions far closer to the quasar sight-line than in most previous studies of high-redshift intervening DLAs, and the two objects we report mark the closest detected high-redshift intervening DLA candidates yet to any quasar sight line. Sections 2, 3, and 4 describe the observations, data reduction, and the subtraction of the quasar point spread functions. Our results are described in section 5. Section 6 describes various tests of our data analysis procedures, carried out to investigate whether the features seen after PSF subtraction are real. A summary of the results of the various data analysis tests is given in subsection 6.12. (Readers interested mainly in the scientific discussion can go directly from section 5 to subsection 6.12.) Finally, sections 7 and 8 discuss the implications of our observations for sizes, environment, and star-formation rates of DLA galaxies. ",
        "conclusions": "With deep diffraction-limited NICMOS images of LBQS 1210+1731, we have probed regions far closer to the quasar sight-line than in most previous studies of high-redshift intervening DLAs. The two objects we report mark the closest detected high-redshift DLA candidates yet to any quasar sight line. Our continuum and H$\\alpha$ images of the $z=1.89$ DLA toward LBQS 1210+1731 suggest that this DLA is not a big galaxy with high SFR, but may be compact (2-3 $h_{70}^{-1}$ kpc in size), probably consisting of multiple sub-units. Assuming no dust extinction of H-$\\alpha$ emission, we place a 3 $\\sigma$  upper limit of 4.0 $h_{70}^{-2}$ M$_{\\odot}$ yr$^{-1}$ on the star formation rate, for $q_{0}=0.5$ . Our continuum and H$\\alpha$ observations are consistent with the hierarchical models, in which DLAs arise in several sub-galactic clumps or dwarf galaxies, which eventually come together to form the present-day galaxies (see, e.g., York et al. 1986; Matteucci et al. 1997). Indeed, theoretical simulations of merging proto-galactic fragments in cold dark matter cosmologies (e.g., Haehnelt et al. 1998), low surface brightness galaxies (e.g., Jimenez et al. 1999), and collapsing halos with merging clouds (e.g., McDonald \\& Miralda-Escud'e 1999) have also been found to reproduce the observed properties of DLAs (asymmetric line profiles of metal absorption lines, metallicities, H I content etc.) The small sizes of high-$z$ DLAs suggested by our observations are also consistent with the small sizes of galaxies seen in other independent high-redshift observations, e.g., in the NICMOS Hubble Deep Field observations (Thompson et al. 1999). Together, these observations may be indications that while star formation had begun long before $z=2$ resulting in some chemical enrichment, most of the dynamical assembly of galaxies as we know them today occurred more recently, and at $z \\sim 2$, the various constituent units were still coming together. However, it cannot be ruled out that the DLA toward LBQS1210+1731 is a large low surface brightness galaxy with a low SFR, which is below our detection limit even in the F160W image.  We point out that our conclusions are, nevertheless, based on detailed observations of only one high-$z$ DLA. It is quite possible that different DLAs have different rates of evolution because of different physical conditions. Indeed, this is suggested by the large scatter in the metallicity-redshift relation of DLAs (see, e.g., Pettini et al. 1999 and references therein). The NICMOS observations of other DLAs from our sample are currently being analyzed and will help to explore the generality of our conclusions. To improve the statistics of the DLA imaging studies, it is necessary to obtain high spatial resolution near-IR images of more high-redshift DLAs. It would be very valuable to carry out a deeper near-IR imaging survey of more DLAs with HST, if the NICMOS cryocooler or the near-IR channel of WFC3 becomes available in the near future. A major advantage of such HST observations will be a relatively stable PSF compared to that currently achieved with any ground-based telescope, which is crucial for the detection of DLAs. It will also be of great interest to complement the HST observations with observations from adaptive optics systems on large ground-based telescopes. Although these systems will not initially have the relatively stable PSF offered by HST, they will be able to achieve even higher spatial resolution and higher imaging sensitivity. Such future space and ground-based observations will provide further insight into the structure and nature of DLA galaxies, and thereby help to constrain theoretical models of the formation and evolution of galaxies."
    },
    "0002/astro-ph0002390_arXiv.txt": {
        "abstract": " ",
        "introduction": "Powerful radio galaxies provide one of the most exciting opportunities to trace the evolution of galaxies and their activity over the history of the Universe. There is a rapid evolution of the power of active galactic nuclei (AGN), such as quasars and radio galaxies, toward $z\\sim3$, and such evolution may reflect some aspects of the phenomena of galaxy formation and evolution.  Powerful radio galaxies have an advantage over quasars in studying the structures in hosts since they are not swamped by the bright nucler light.  One of the most striking properties of high-redshift powerful radio galaxies (HzPRGs) is the `alignment effect' observed in their optical images (McCarthy et al. 1987; Chambers et al. 1987; Rigler et al. 1992). Rest-frame ultraviolet (UV) radiation tends to have an elongated structure aligned with the radio-jet axis. The origin of this aligned UV radiation is still not clear; while the scattered non-thermal radiation contributes to the observed UV flux, at least to some extent, in many radio galaxies, the stellar photospheric emission and/or the nebular continuum emission has also been observed in some sources (e.g., Cimatti et al. 1996). In any case, UV emission is very likely to be the {\\it consequence} of nuclear activity, and near-infrared observations may be more essential to reveal the true host-galaxy structure.  3C 324 at $z=1.206$ has been extensively studied as one of the `proto-typical' objects which show the optical alignment effect. Its total radio power at 5 GHz is large, log $P_{\\rm 5 GHz}$= 27.75 ($H_0$ = 50 km s$^{-1}$Mpc$^{-1}$, $q_0$ = 0.5; hereafter we use this set of cosmological parameters unless noted), which implies the existence of a very powerful AGN. Indeed, Cimatti et al. (1996) detected a broad Mg {\\sc II} $\\lambda$ 2800 \\AA\\ emission line in the polarized spectrum, and thus established that 3C 324 hosts a quasar obscured by dust. The radio source has two distinct radio lobes with a separation of $\\sim 11$$^{\\prime\\prime}$ (95 kpc at $z=1.206$; Fernini et al. 1993; Best et al. 1998b) at a position angle of 71$^\\circ$ (Dunlop, Peacock 1993). Best et al. (1998b) detected a plausible radio core at frequencies of 8.2 and 4.7 GHz. Hubble Space Telescope (HST) WFPC2 observations with the ${\\rm F702W}$ and ${\\rm F791W}$ filters (Longair et al. 1995; Dickinson et al. 1995; Best et al. 1997, 1998a) revealed in detail the elongated patchy structure of the rest-frame UV emission, which is collinear with the radio axis. The radio core is located in a gap between the two central optical knots (Best et al. 1998b), which suggests that this gap is due to a dust lane which obscures the active nucleus. The existence of a significant amount of dust is supported by the detection of sub-mm radiation which is not associated with the synchrotron radiation (Best et al. 1998c).  At near-infrared (NIR) wavelengths, previous ground-based observations under $\\sim$ 1$^{\\prime\\prime}$ seeing conditions have detected the red host galaxy (Dunlop, Peacock 1993; Dickinson et al. 1995;  Best et al. 1997, 1998a). Dunlop and Peacock (1993) showed that the host galaxy has a somewhat elongated structure whose position angle (PA) is $\\sim 75^\\circ$, closely aligned with the radio axis. Best et al. (1998a) argued that the galaxy is as bright as $K=16.99$ mag (9$^{\\prime\\prime}$ aperture) and  has a light profile well represented by a de Vaucouleurs law with an effective radius of 2$^{\\prime\\prime}$.2 (19 kpc). Dickinson et al. (1995) pointed out that the $K$-band light peak coinsides with the gap of the optical knots. It has been known that 3C 324 is located in a galaxy cluster (Dickinson et al. 1997; Kajisawa et al. 1999 and references therein) and is the brightest galaxy in the cluster at z = 1.21.  We have observed 3C 324 with the Subaru telescope equipped with the Cooled Infrared Spectrograph and Camera for OHS (CISCO: Motohara et al. 1998), which provides a $\\sim 2$$^{\\prime\\prime}$ $\\times 2$$^{\\prime\\prime}$  field of view at a sampling of 0$^{\\prime\\prime}$.116  pixel$^{-1}$. The observations were made during the telescope commissioning period and good image quality with 0$^{\\prime\\prime}$.3--0$^{\\prime\\prime}$.4 seeing (FWHM of stellar images) was achieved. In this paper, we investigate the NIR properties of the host galaxy of 3C 324. Thanks to the high spatial resolution, we can directly compare the light distribution in the optical images obtained with the deep HST observations without seriously degrading the HST data. In section 2, we briefly describe the observations and the data reduction. We show the light distribution and color map of the host galaxy in section 3 and discuss its properties as a brightest cluster galaxy in section 4. We present our conclusions in section 5. ",
        "conclusions": "The host galaxy of the powerful radio galaxy 3C 324 was observed with the Subaru telescope under good seeing conditions. The host galaxy is clearly resolved and seen to be a spheroidal galaxy well approximated by a de Vaucouleurs profile. The effective (half-light) radius evaluated from profile fitting is 1$^{\\prime\\prime}$.3 (11.2 kpc), which is about half the value previously published in the literature, while a curve of growth analysis produces a value of 0$^{\\prime\\prime}$.8. After subtraction of the model galaxies, we clearly detect the `aligned component' in the $K^\\prime$-band image. The disagreement between the effective radius obtained in the profile fitting and the half-light radius in the growth-curve analysis may be due to this `aligned component' and/or to a contribution from the obscured AGN.  The peak of the $K^\\prime$-band light coincides with the position of the radio core, which strongly implies that the engine of the powerful radio sources is indeed hosted at the nucleus of the giant elliptical galaxy. The NIR peak also corresponds to the gap in the rest-frame UV emission, which may be due to a dust lane. It is very likely that we see the obscuring structure from an almost edge-on view.  The host galaxy has a very red $R_{{\\rm F702W}}-K$ color and the near-infrared light of the galaxy is likely to be dominated by an old stellar population, while the relatively blue $B_{{\\rm F450W}}-R_{{\\rm F702W}}$ color suggests that there may be some small amount of star-formation activity.  The colors of the 'aligned' components located inside the host galaxy, which are obtained after subtracting the host component, may be explained by nebular continuum emission with a small amount of a dust while those outside the host galaxy are better modeled by optically-thin dust scattering of the nuclear light.   \\vspace{0.5cm}  The authors are indebted to all members of the Subaru Observatory, NAOJ, Japan. We thank Nobuo Arimoto for kindly providing the Kodama and Arimoto evolutionary synthesis models. TY thanks Takashi Murayama for useful discussions.  We thank Dr. Marc Dickinson, the referee, for the invaluable comments. This research was supported by grants-in-aid for scientific research of the Ministry of Education, Science, Sports and Culture (08740181, 09740168). This work was also supported by the Foundation for the Promotion of Astronomy of Japan.  This work is based in part on observations with the NASA/ESA Hubble Space Telescope, obtained from the data archive at the Space Telescope Science Institute, U.S.A., which is operated by AURA, Inc.\\ under NASA contract NAS5--26555. The Image Reduction and Analysis Facility (IRAF) used in this paper is distributed by National Optical Astronomy Observatories. U.S.A., operated by the Association of Universities for Research in Astronomy, Inc., under contact to the U.S.A. National Science Foundation. \\clearpage"
    },
    "0002/astro-ph0002029_arXiv.txt": {
        "abstract": "A problem of mass flow in the immediate vicinity of a planet embedded in a protoplanetary disk is studied numerically in two dimensions. Large differences in temporal and spatial scales involved suggest that a specialized discretization method for solution of hydrodynamical equations may offer great savings in computational resources, and can make extensive parameter studies feasible. Preliminary results obtained with help of Adaptive Mesh Refinement technique and high-order explicit Eulerian solver are presented. This combination of numerical techniques appears to be an excellent tool which allows for direct simulations of mass flow in vicinity of the accretor at moderate computational cost. In particular, it is possible to resolve the surface of the planet and to model the process of planet growth with minimal set of assumptions. Some issues related to visualization of the results and future prospects are discussed briefly. ",
        "introduction": "Extremely small temporal and spatial scales involved in the problem of accretion onto a protoplanet necessitate the use of nonuniform discretization in the vicinity of the accretor. In our study we used adaptive mesh refinement (AMR) method combined with a high-resolution Godunov-type advection scheme ({\\sc amra}, Plewa \\& M\\\"uller 2000). The AMR discretization scheme follows the approach of Berger and Colella (1989). The computational domain is covered by a set of completely nested {\\em patches} occupying {\\em levels}. The levels create a refinement hierarchy. As one moves toward higher levels, the numerical resolution increases by a prescribed integer factor (separate for every direction). The net flow of material between patches at different levels is carefully accounted for in order to preserve conservation properties of hydrodynamical equations. Boundary data for child patches are either obtained by parabolic two-dimensional conservative interpolation of parental data or set according to prescribed boundary conditions.  Hydrodynamical equations are solved with the help of the Direct Eulerian Piecewise-Parabolic Method (PPMDE) of Colella \\& Woodward (1984), as implemented in {\\sc herakles} solver (Plewa \\& M\\\"uller 2000). Simulations have been done in spherical polar coordinates in a frame of reference corotating with the protoplanet. {\\sc herakles} guarantees exact conservation of angular momentum which is particularly important in numerical modeling of disk accretion problems. The use of its multifluid option with tracer materials distributed within disk (not presented here) allows to identify the origin of the material accreted onto protoplanet. The {\\sc amra} code is written purely in FORTRAN 77 and has been successfully used on both vector supercomputers and superscalar cache-oriented workstations. Its parallelization on shared memory machines exploits microtasking (through the use of vendor-specific directives) or the OpenMP standard. ",
        "conclusions": ""
    },
    "0002/astro-ph0002503_arXiv.txt": {
        "abstract": "We employed recently computed evolutionary white-dwarf models with helium cores, supplemented by heavier models with carbon-oxygen cores, in order to investigate the ages of millisecond pulsar systems based on the cooling properties of the compact companions. Contrary to the behaviour of more massive white dwarfs, the evolutionary speed  of low-mass white-dwarf models is substantially slowed down by ongoing hydrogen burning. By comparing the cooling ages of these models with the spin-down ages of the pulsars for those systems for which reasonable information about the compact companions is available, we found good correspondence between both ages. Based on these models any revisions concerning the temporal evolution of millisecond pulsars do not appear to be necessary. ",
        "introduction": "Millisecond pulsars are thought to be components of low-mass binary systems in their final stage of evolution: the neutron star which has been spun up by accretion of matter from a low-mass evolved companion is now being slowed down by emission of magnetic dipole radiation (recycled radio pulsar). The companion, after having transferred most of its envelope mass towards the neutron star, remains as a white dwarf of rather a low mass whose core consists, in the majority of the known cases, of helium. The characteristic age, or so-called {\\em spin-down age}, of the (recycled) pulsar depends on the physics how the neutron star's rotational energy  is converted into non-thermal emission of electromagnetic energy. On the other hand, the white-dwarf age is ruled by the white dwarf's thermo-mechanical structure and the transformation of gravothermal energy content into thermal emission of photons from the surface. Any age determinations of the pulsar and the dwarf component should give the same answer, provided our physical understanding of the pulsar's slow-down processes and the white dwarf's cooling properties  is correct.  So far, no general consensus on this matter has been achieved. Under the assumption that the cooling properties of low-mass white dwarfs are ruled by rather simple laws as is known from evolutionary calculations of more massive white dwarfs with carbon-oxygen cores (cf.\\ Iben \\& Tutukov \\cite{IT84}; Koester \\& Sch\\\"onberner \\cite{KS86}, Bl\\\"ocker \\cite{BL95}), large age differences between the pulsars and their dwarf companions have been found. In general, the white dwarfs appear to be much younger than the pulsars (cf.\\ Hansen \\& Phinney \\cite{HP98b} for a recent, detailed account). The best-studied example is the \\object{PSR J1012+5307} system, for which Lorimer et al.\\ (\\cite{LFLN95}) determined 7 Gyr for the spin-down age of the pulsar, but only about 0.3 Gyr for the white dwarf's age. Note that the usual spin-down age determinations are based on the assumption that the initial rotational period after completion of the spin-up by accretion is much smaller then the present one, and that the pulsar emits magnetic dipole radiation (braking index $n=3$). A summary of the assumptions inherent in the derivation of characteristic or spin-down ages of pulsars is given in Hansen \\& Phinney (\\cite{HP98b}). A discrepant result as found for \\object{PSR J1012+5307}, if true, would have important consequences for the details of the accretion process and the following spin-down phase (cf.\\ Burderi et al.\\ \\cite{BKW96}).  A larger sample of millisecond pulsar systems with white-dwarf companions has recently been investigated by Hansen \\& Phinney (\\cite{HP98b}), using a grid of low-mass white-dwarf sequences especially computed for this purpurse (Hansen \\& Phinney \\cite{HP98a}). In most cases spin-down and cooling ages appeared to be discrepant to various degrees, and the authors were able to constrain the initial spin periods and spin-up histories for individual systems, especially also for the \\object{PSR J1012+5307} system. However, the white-dwarf models which this study is based on, are generated from ad-hoc assumed initial configurations. These configurations appear not to be consistent with respect to the thermo-mechanical structures and unprocessed, hydrogen-rich envelopes with what would be adequate for companions in these pulsar binary systems.  The early investigations concerning the evolution of helium white dwarfs made by Webbink (\\cite{W75}) indicated that the final cooling is slowed down considerably by ongoing hydrogen burning via the pp cycle. Obviously the cooling behaviour of low-mass white dwarfs depends on the size of the still unprocessed hydrogen-rich envelope, i.e.\\ whether this envelope is massive enough as to sustain burning temperatures at its bottom for a long time span. The Webbink (\\cite{W75}) white-dwarf models are, however, just evolved main sequence stars without any consideration of mass loss.  Since white-dwarf envelope masses cannot be guessed from first principles, they must rather be determined by detailed evolutionary calculations. A step in this direction was made by Alberts et al.\\ (\\cite{ASH96}) and Sarna et al.\\ (\\cite{SAAM98}) who modelled % the \\object{PSR J1012+5307} system and in particular the evolution of the mass giving companion. % It turned  out that the donor shrinks below its Roche lobe while still having a rather massive hydrogen-rich envelope which is able to keep hydrogen burning dominant even through the white-dwarf cooling phase. The evolution was slowed down to such an extent that the discrepancy with the spin-down age of the pulsar vanished completely.  Strictly speaking the strength of hydrogen burning, and hence the cooling age of an observed white dwarf, depends on the size of the envelope before entering the cooling path. This envelope mass can be reduced because of thermal instabilities of the burning shell when the CNO rate dies out, namely by \\begin{itemize} \\item enhanced hydrogen consumption during the instability (flash) itself, and by \\item a possible Roche-lobe overflow driven by the rapid envelope expansion. \\end{itemize} The latter case was dominant for the evolution of the Iben \\& Tutukov (\\cite{IT86}) 0.3\\,M$_{\\sun}$ helium white-dwarf model: Roche-lobe overflow due to the flash-driven envelope expansions reduced the envelope mass {\\em below\\/} the critical value necessary for hydrogen burning. The white-dwarf models of Webbink (\\cite{W75}) and Sarna et al.\\ (\\cite{SAAM98}) experienced phases of unstable hydrogen burning for \\mbox{$ M \\la 0.2$~M$_{\\sun}$} (but see Driebe et al.\\ \\cite{DBSH99} for a discussion).  Recently Driebe et al.\\ (\\cite{DSBH98}) published a grid of evolutionary tracks for helium white-dwarf models which were generated by enhanced mass loss applied at different positions along the red-giant branch of a 1\\,M$_{\\sun}$ sequence (see also Iben \\& Tutukov \\cite{IT86}, Castellani et al.\\ 1994). This method mimicks to some extent the mass transfer in  binary systems and allows to get reliable post-red-giant configurations which are very useful for the interpretation of observations. Driebe et al.\\ (\\cite{DSBH98}) covered the whole mass range of interest, and they demonstrated that \\begin{itemize} \\item the anti-correlation between core mass and size of envelope (cf.\\ Bl\\\"ocker et al.\\ \\cite{BHDBS97}) determines later the nuclear activity along the cooling branch, and that \\item thermal instabilities of the hydrogen-burning shell appear to be restricted to the mass range of approximately 0.2 to 0.3~M$_{\\sun}$. % \\end{itemize} The absence of thermal flashes below $ M= 0.2$~M$_{\\sun}$ agrees well with the results of Alberts et al.\\ (\\cite{ASH96}) but disagrees with those of Sarna et al.\\ (\\cite{SAAM98}). Nevertheless, the cooling times of our models are in excellent agreement with both studies. From the given parameters of the white-dwarf component in the \\object{PSR J1012+5307} system, Driebe et al.\\  (\\cite{DSBH98}) determined then its age to be of $ 6 \\pm 1$~Gyr, in good agreement with the pulsar's spin-down age of $ 7.0 \\pm 1.4$~Gyr (Lorimer et al.\\ \\cite{LFLN95}).  The latest effort in a better understanding of the combined pulsar-white dwarf systems is that of Burderi et al.\\ (\\cite{BKW98}). They took the pulsar spin-down ages at their face value and concluded that the standard assumption for the white-dwarf cooling (i.e.\\ without nuclear burning) complies with the observations, except for masses below approx.\\ 0.2~M$_{\\sun}$. There are, however, some facts that we would like to point out: Burderi et al.\\  (\\cite{BKW98}) used data 'renormalized' to a standard luminosity of $ 10^{-2}$~L$_{\\sun}$, whereby it remains unclear how ages can be renormalized if the temporal evolution of the systems is not known a priori. Furthermore, they extrapolated existing white-dwarf cooling models into mass regimes where they are not valid anymore.  Because of its importance we  felt the necessity to reconsider the whole issue by utilizing more realistic evolutionary models for low-mass white dwarfs. We will show in the next section that with such models a consistent description of those millisecond pulsar binary systems can be achieved for which sufficiently accurate data is available. ",
        "conclusions": "From the present paper, together with previous efforts which concentrated solely on the \\object{PSR J1012+5307} system (Alberts et al.\\ \\cite{ASH96}; Sarna et al.\\ \\cite{SAAM98}; Driebe et al.\\ \\cite{DSBH98}), it becomes obvious that a consistent description of the millisecond binary systems with compact companions can only be achieved by using % evolutionary model calculations of white dwarfs which include their complete pre-white-dwarf history. The key to the solution of the apparent age paradoxon between the pulsar and its white-dwarf companion is the fact that low-mass white dwarfs have massive, still unburnt envelopes that sustain hydrogen burning at their bases for a long time. Hydrogen burning slows down the cooling of a low-mass white dwarf to such an extent that cooling ages become comparable to, or may even exceed, observed pulsar spin-down ages.  Employing our evolutionary helium white-dwarf models, supplemented by those with carbon-oxygen cores, we demonstrated that, next to the already studied PSR J1012+5307 system, also in other millisecond pulsar binary systems with reliable information on pulsar age and companion properties, as mass and temperature, reasonable agreement between the components' ages is achieved. The use of white-dwarf models with ad hoc assumed envelope masses may lead to erroneous interpretations since these envelope masses are usually much smaller than those which follow from complete evolutionary calculations.  It appears to us that upon using  realistic white-dwarf models in interpreting millisecond binary systems there is no need to modify existing ideas of the spin-down process of pulsars. Clearly a larger sample of well-studied systems like \\object{PSR J1012+5307} would be very important in investigating more precisely the cooling theory of white dwarfs {\\it and\\/} the braking of radio pulsars."
    },
    "0002/astro-ph0002053_arXiv.txt": {
        "abstract": "The theory of coronal evaporation predicts the formation of an inner hole in the cool thin accretion disk for mass accretion rates below a certain value ($\\approx$ 1/50 of the Eddington mass accretion rate) and the sudden disappearance of this hole when the mass accretion rate rises above that value. The inner edge of the standard thin disk then suddenly shifts inward vvvvfrom about a few hundred Schwarzschild radii to the last stable orbit. This appears to quantitatively account for the observed transitions between  hard and soft spectral states at critical luminosities. Due to the evaporation process the matter accreting in the geometrically thin disk changes to a hot coronal flow which proceeds towards the black hole as an advection-dominated accretion flow (ADAF; for a review see Narayan et al. 1998). ",
        "introduction": "For a decade it has been known that the spectra of X-ray novae show changes from a soft state at high luminosity to a hard state when the luminosity has declined during the outburst (Tanaka 1989). The persistent canonical black hole system Cyg X-1 also undergoes occasional transitions between its standard low luminosity (hard) state and a soft state (see Fig. 1). Such changes between the two spectral states have been observed for several systems, regardless of whether the compact object is a neutron star (Aql X-1, 1608-522) or a black hole (GS/GRS 1124-684, GX 339-4) (Tanaka \\& Shibazaki 1996). Here we concentrate on black hole sources. Observations show that the phenomenon always occurs at a luminosity around $10^{37}\\rm{erg/s}$, which corresponds to a mass accretion rate of about $10^{17}\\rm{g/s}$ (Tanaka 1999).  The two spectral states are thought to be related to different states of accretion: (1) the soft spectrum originates from a thin disk which extends down to the last stable orbit plus a corona above the disk, (2) the hard spectrum originates from a thin disk outside a transition radius $r_{tr}$ and a coronal flow/ ADAF inside. The spectral transitions of Nova Muscae 1991 and Cygnus X-1 were modelled based on this picture by Esin et al. (1997, 1998). The value of $r_{tr}$ was taken as the maximal distance $r$ for which an ADAF with that accretion rate can exist (``strong ADAF proposal\", Narayan \\& Yi 1995). We determine the location of the inner edge of the thin disk from the equilibrium between it and the corona above. ",
        "conclusions": "We understand the spectral transition as related to a critical mass accretion rate. For rates ${\\dot M} \\ge {\\dot M_{\\rm{crit}}}$ (the peak coronal mass flow rate) the standard disk reaches inward to the last stable orbit and the spectrum is soft. Otherwise the ADAF in the inner accretion region provides a hard spectrum. At ${\\dot M_{\\rm{crit}}}$ the transition between dominant advective losses further out and dominant radiative losses further in occurs. Except for the difference between the sub-virial temperature of the corona and the closer-to-virial temperature of an ADAF of the same mass flow rate, this same critical radius is predicted by the ``strong ADAF proposal\" (Narayan \\& Yi (1995). In general however, the strong ADAF proposal results in an ADAF region larger than that which the evaporation model yields.   The transition between the two spectral states  has been observed for black hole and neutron star systems, in persistent and transient sources (Tanaka \\& Shibazaki 1996, Campana et al.1998). This points to similar physical accretion processes. Menou et al. (1999) already discussed the accretion via an ADAF in neutron star transient sources. Our results should also be applicable.  The relations for a 6 $M_\\odot$ black hole plotted in Fig. 2 can be scaled to other masses: in units of Schwarzschild radii and Eddington accretion rates the plot is universal. The application to disks around supermassive black holes implies interesting conclusions for AGN."
    },
    "0002/astro-ph0002265_arXiv.txt": {
        "abstract": "\\rightskip 0pt \\pretolerance=100 \\noindent We present the results of hard-X-ray observations of four broad-line radio galaxies (BLRGs) with the {\\it Rossi X-Ray Timing Explorer} ({\\it RXTE}). The original motivation behind the observations was to search for systematic differences between the BLRGs and their radio-quiet counterparts, the Seyfert galaxies. We do, indeed, find that the Fe~K$\\alpha$ lines and Compton ``reflection'' components, which are hallmarks of the X-ray spectra of Seyferts galaxies, are weaker in BLRGs by about a factor of 2. This observational result is in agreement with the conclusions of other recent studies of these objects. We examine several possible explanations for this systematic difference, including beaming of the primary X-rays away from the accretion disk, a low iron abundance, a small solid angle subtended by the disk to the primary X-ray source, and dilution of the observed spectrum by beamed X-rays from the jet. We find that a small solid angle subtended by the disk to the primary X-ray source is a viable and appealing explanation, while all others suffer from drawbacks. We interpret this as an indication of  a difference in the inner accretion disk structure between Seyfert galaxies  and BLRGs, namely that the inner accretion disks of BLRGs have the form of an ion-supported torus or an advection-dominated accretion flow, which irradiates the geometrically thin outer disk. ",
        "introduction": "An important issue in our study of active galactic nuclei (hereafter AGNs) is the as yet unexplained difference between radio loud and radio-quiet objects. All AGNs are thought to be powered by accretion of matter onto a supermassive black hole, presumably via an equatorial accretion disk. From a theoretical perspective, the accretion disk is an essential ingredient for the formation of radio jets, although the exact mechanism is not well known (see the review by Livio 1996). The observational evidence\\footnote{The most direct observational evidence for the presence of accretion disks in AGNs takes the form of Fe~K$\\alpha$ lines with disk-like profiles in the X-ray spectra of Seyfert galaxies (e.g., Tanaka et al. 1995; Nandra et al. 1997b) and double-peaked H$\\alpha$ lines in the optical spectra of BLRGs (Eracleous \\& Halpern 1994).} suggests that both radio-loud and radio-quiet AGNs harbor accretion disks, but it is a mystery why well-collimated, powerful, relativistic radio jets only exist in the former class of object. The origin of the difference could lie in the nature of the host galaxy. At low redshifts ($z<0.5$) radio-loud AGNs are found only in elliptical galaxies, whereas radio-quiet AGNs can have either elliptical or spiral hosts (Smith et al. 1986; Hutchings, Janson, \\& Neff 1989; V\\'eron-Cetty \\& Woltjer 1990; Dunlop et al. 1993; Bahcall et al. 1997; Boyce et al. 1998). This observational trend has led to the suggestion that the interstellar medium of the host galaxy may play an important role in the propagation and collimation of the radio jets on large scales (e.g., Blandford \\& Levinson 1995; Fabian \\& Rees 1995). Alternatively, one may seek the fundamental cause of the difference between radio-loud and radio-quiet AGNs in the properties of their accretion flows or the properties of their central black holes. One possibility is that radio-loud AGNs may harbor rapidly spinning black holes, whose energy is extracted electromagnetically via the Blandford \\& Znajek (1977) mechanism and used to power the radio jets. Rapidly spinning black holes could be associated with elliptical galaxies if both the black hole and the host galaxy result from the merger of two parent galaxies, each with its own nuclear black hole (see Wilson \\& Colbert 1995).  Another possibility is that the inner accretion disks of radio-quiet AGNs are geometrically thin and optically thick throughout (Shakura \\& Sunyaev 1973) while the inner accretion disks of  radio-loud AGNs are ion-supported tori (Rees et al. 1982; known today as advection-dominated accretion flows, or ADAFs, after the work of Narayan \\& Yi 1994, 1995). Because ADAFs are nearly spherical and parts of the flow are unbound, they can lead to the formation of outflows (e.g., Blandford \\& Begelman 1999). In either of the  above pictures, an additional mechanism may be necessary to {\\it collimate} the radio jets (see for example, Blandford \\& Payne 1982; Meier 1999).  If the difference between radio-loud and radio-quiet AGNs is related to differences in their central engines, it would lead to observable differences in their respective X-ray spectra, in particular in the properties (profiles and equivalent width) of the Fe~K$\\alpha$ lines and in the shape of the continuum. This is because the Fe~K$\\alpha$ lines are thought to result from fluorescence of the dense gas in the geometrically thin and optically thick regions of the disk (e.g., George \\& Fabian 1991; Matt et al 1992). Similarly, the continuum above 10~keV is thought to include a significant contribution from X-ray photons from  the ``primary'' X-ray source, near the center of the disk, which undergo Compton scattering (``reflection'') in the same regions of the disk where the  Fe~K$\\alpha$ line is produced (e.g., Lightman \\& White 1988; George \\& Fabian 1991; Matt, Perola, \\& Piro 1991). It is, therefore, extremely interesting that studies of the X-ray spectra of broad-line radio galaxies (BLRGs) by Zdziarski et al. (1995) and Wo\\'zniak et al. (1998) found them to be systematically different from those of (radio-quiet) Seyfert  galaxies. In particular, these authors found that the signature of Compton reflection, which is very prominent in the spectra of Seyfert galaxies above 10~keV (Pounds et al. 1989; Nandra \\& Pounds 1994), is weak of absent in the spectra of BLRGs. Moreover, Wo\\'zniak et al. (1997) found that the Fe~K$\\alpha$ lines of BLRGs are  narrower and weaker than those of Seyfert galaxies.  Motivated by the above theoretical considerations and observational results, we have undertaken a systematic study of the X-ray spectra of radio-loud AGNs with {\\it ASCA} and {\\it RXTE} in order to characterize their properties. Our  main goal is to compare their spectroscopic properties with those of Seyfert galaxies and test the above ideas for the origin of the difference between the two classes. In our re-analysis of archival {\\it ASCA} spectra of BLRGs and radio-loud quasars (Sambruna, Eracleous, \\& Mushotzky 1999)  we found that the Fe~K$\\alpha$ lines of some objects are indeed weaker and narrower than those of Seyferts, in agreement with the findings of Wo\\'zniak et al. (1997). In other objects, however, the uncertainties are large enough that we cannot reach firm conclusions, thus we have not been able to confirm this result in general. In this paper we present the results of new observations of four BLRGs with {\\it RXTE}, aimed at measuring the shape of their hard X-ray continuum and the equivalent width of their Fe~K$\\alpha$ lines. As such, these observations complement our study of the {\\it ASCA} spectra of these objects. In \\S2 we describe the observations and data screening. In \\S3 we present and discuss the light curves and in \\S4 we compare the observed spectra with models. In \\S5 we discuss the implications of the results, while in \\S6 we summarize our conclusions.  Throughout this paper we assume a Hubble constant of $H_0=50~{\\rm km~s^{-1}~Mpc^{-1}}$ and a deceleration parameter of $q_0=0.5$.   \\begin{deluxetable}{lcrcccc} \\tablenum{1} \\tablewidth{6.5in} \\tablecolumns{6} \\tablecaption{Target Objects, Basic Properties, and Observation Log} \\tablehead{ & & & \\multicolumn{2}{c}{$N_{\\rm H}$ (cm$^{-2}$)} & \\nl & & & \\multicolumn{2}{l}{\\hrulefill} & \\colhead{Observation} & Duration \\nl \\colhead{Object}  & \\colhead{$z$}  & \\colhead{$i$\\tablenotemark{\\;a}}  & \\colhead{Galactic\\tablenotemark{\\;b}}  & \\colhead{{\\it ASCA}\\tablenotemark{\\;c}} & \\colhead{Start Time (UT)} & \\colhead{(hours)} } \\startdata 3C 111    & 0.048 & $37^{\\circ}>i>24^{\\circ}$           & $1.2\\times 10^{22}$ & $9.63\\times 10^{21}$ & 1997/3/22 01:23 & 62 \\nl 3C 120    & 0.033 & $14^{\\circ}>i>1^{\\circ}\\phantom{4}$ & $1.2\\times 10^{21}$ & $1.65\\times 10^{21}$ & 1998/2/13 04:53 & 58 \\nl Pictor A  & 0.035 & $i>24^{\\circ}$                      & $4.2\\times 10^{20}$ & $8.30\\times 10^{20}$ & 1997/5/08 02:21 & 82 \\nl 3C 382    & 0.057 & $i>15^{\\circ}$                      & $6.7\\times 10^{20}$ & $6.70\\times 10^{20}$ & 1997/3/28 23:26 & 47 \\nl \\tablenotetext{a\\;} {The inclination angle of the radio jet (see Eracleous \\& Halpern 1998, and references therein).} \\tablenotetext{b\\;} {The Galactic equivalent H{\\sc\\, i} column density. {\\it References}: 3C~111: Bania, Marscher, \\& Barvainis (1991); 3C~120: Elvis et al. (1989); Pictor~A: Heiles \\& Cleary (1979); 3C 382: Murphy et al. (1996). } \\tablenotetext{c\\;} {The Galactic equivalent H{\\sc\\, i} column density as measured by the {\\it ASCA} SIS (taken from Sambruna et al. 1999).} \\enddata \\end{deluxetable}   \\begin{deluxetable}{lcccccc} \\tablenum{2} \\tablewidth{6.5in} \\tablecolumns{7} \\tablecaption{Exposure Times and Count Rates} \\tablehead{ & \\multicolumn{3}{c}{PCA (2.5--30 keV)} & \\multicolumn{3}{c}{HEXTE cluster 0 (20-100 keV)}\\nl & \\multicolumn{3}{c}{\\hrulefill} & \\multicolumn{3}{c}{\\hrulefill}\\nl \\colhead{} & \\colhead{Exp.} & \\colhead{Source} & \\colhead{Backg.} & \\colhead{Exp.} & \\colhead{Source} & \\colhead{Backg.} \\nl \\colhead{} & \\colhead{Time} & \\colhead{Count Rate} & \\colhead{Count Rate} & \\colhead{Time} & \\colhead{Count Rate} & \\colhead{Count Rate} \\nl \\colhead{Object} & \\colhead{(s)} & \\colhead{(s$^{-1}$ PCU$^{-1}$)} & \\colhead{(s$^{-1}$ PCU$^{-1}$)} & \\colhead{(s)} & \\colhead{(s$^{-1}$)} & \\colhead{(s$^{-1}$)} } \\startdata 3C 111    & 33,440 & 22.9 & 39.1 & 12,748 & 7.5 & 158.4 \\nl 3C 120    & 55,904\\tablenotemark{\\;a} & 25.2 & 39.9 & 17,572 & 3.7 & 187.5 \\nl Pictor A  & 30,240\\tablenotemark{\\;a} & 9.6 & 38.7 & 12,461 & 9.4 & 166.9 \\nl 3C 382    & 28,192 & 13.9 & 22.0 & 10,606 & 3.0 & 114.9 \\nl \\tablenotetext{a\\;} {The PCA exposure time refers to PCUs 0, 1, and 2. PCUs 3 and 4 were off during part of the observation.} \\enddata \\end{deluxetable} ",
        "conclusions": "Our study of the hard X-ray spectra of 4 BLRGs observed with {\\it RXTE} has shown them to be systematically different from those of Seyfert galaxies. In particular, the Fe~K$\\alpha$ lines and Compton reflection humps in the spectra of BLRGs are at least a factor of 2 weaker than those observed in the spectra of Seyfert galaxies. This result is consistent with the conclusions of previous studies and is supported by the results of {\\it SAX} observations of  the same targets. After examining several possible explanations for this difference, we conclude that the most likely one is that the solid angle subtended by the source of these spectral features to the primary X-ray source is a factor of 2 smaller in BLRGs than in Seyferts. Since this reprocessing medium is thought to be the accretion disk, we interpret this difference as the result of a difference in the accretion disk structure between BLRGs and Seyferts. More specifically, we argue that if the inner accretion disks of BLRGs have the form of an ion-supported torus (or an ADAF) that irradiates the outer disk, then the observed differences can be explained. We find this explanation particularly appealing because ADAFs offer a possible way of producing the radio jets in these objects and because such a disk structure can also account for the double-peaked Balmer line profiles observed in these objects.  This interpretation, as well as some of the other scenarios we have considered, can be tested further by studying the profiles of of the Fe~K$\\alpha$ lines. Unfortunately, such a test has not been possible so far using spectra from {\\it ASCA}, because of their low signal-to-noise ratio. It will be possible, however, to carry out the test using spectra from upcoming observatories such as {\\it XMM} and {\\it Astro-E}. Correlated variations of the intensity of the Fe~K$\\alpha$ line and the strength of the X-ray continuum afford an additional observational test of these ideas. The light travel time between the ion torus and the outer accretion disk is  $\\tau_{\\ell} \\approx 1.7\\,(R/300\\,R_{\\rm g})\\,(M/10^8\\,M_{\\odot})$~days, which means that one may expect lags of this order between line and continuum variations. If, on the other hand, the line is produced very close to the center of the disk, the lags should be considerably shorter, while if the line is produced in the obscuring torus, the lags should be on the order of several years. Monitoring campaigns with {\\it RXTE} may provide the data needed for such a test."
    },
    "0002/astro-ph0002115_arXiv.txt": {
        "abstract": "We report the most recent progress in understanding the emission properties of millisecond pulsars. ",
        "introduction": "} \\setcounter{footnote}{1} \\footnotetext{If two do the same, it is not the same.}   Through intensive research for almost two decades, it has been well established, both in theory and observation, that millisecond pulsars (MSPs) are the end product of mass accretion in binary systems. As MSPs emerge in the radio universe having been given a second chance in life, they are surrounded by magnetospheres which are several orders of magnitude more compact than those of slower rotating pulsars. Inferred magnetic fields close to the surface of MSPs are 3 to 4 orders of magnitude weaker than in normal pulsars while charges at these regions experience an accelerating potential similar to that of normal pulsars.  The impact of the different environment on the emission process in MSP magnetospheres has been a question addressed already shortly after the discovery of a first few such sources.  With the plethora of MSPs detected over the years, a significant sample became available to us, enabling a better understanding of not only MSPs (as radio sources and tools) but slower rotating (normal) pulsars as well. In the following, we will concentrate on {\\em recent} progress, referring to Kramer et al.~(1998, Paper I) on spectra, pulse shapes and beaming fraction; Xilouris et al.~(1998, Paper II) on polarimetry of 24 MSPs; Sallmen (1998) and Stairs et al.~(1999) on multi-frequency polarimetry; Toscano et al.~(1998) on spectra of Southern MSPs; Kramer et al.~(1999b, Paper III) on multi-frequency evolution; and Kramer et al.~(1999a, Paper IV) on profile instabilities of MSPs; but see also the following contributions by Kuzmin \\& Losovsky and Soglasnov.  \\vspace{-0.3cm} ",
        "conclusions": "While we have had to be necessarily brief in reviewing MSP properties, we direct the interested reader to the extensive studies of MSPs presented in the quoted literature. We summarize here our point of view: MSPs emit their radio emission by the same mechanism as normal pulsars.  Some distinct differences may originate from the way they were formed, but most observed features can be explained by very compact magnetospheres. Our data can be explained without any need to invoke deviations from dipolar field lines, although a large number of open questions remain. We need more polarization information at higher frequencies and, in particular, single pulse studies.  These will allow us to study the formation of the profile and its stability, to see if the additional pulse features are distinct from the main pulse, and how the polarization modes behave under the magnifying glass of the blown-up MSP profiles. There are exciting years to come!"
    },
    "0002/astro-ph0002233_arXiv.txt": {
        "abstract": "We combine new CCD {\\it UBV} photometry and spectroscopy with that from the literature to investigate 19 Magellanic Cloud OB associations that contain Wolf-Rayet (WR) and other types of evolved massive stars. Our spectroscopy reveals a wealth of newly identified interesting objects, including early O-type supergiants, a high mass double-lined binary in the SMC, and, in the LMC, a newly confirmed LBV (R~85), a newly discovered WR star (Sk$-69^\\circ$194), and a newly found luminous B[e] star (LH85-10). We use these data to provide precise reddening determinations and construct physical H-R diagrams for the associations. We find that about half of the associations may be highly coeval, with the massive stars having formed over a short period ($\\Delta \\tau <$ 1~Myr).  The (initial) masses of the highest mass {\\it unevolved} stars in the coeval clusters may be used to estimate the masses of the progenitors of WR and other evolved stars found in these clusters. Similarly the bolometric luminosities of the highest mass unevolved stars can be used to determine the bolometric corrections for the evolved stars, providing a valuable observational basis for evaluating recent models of these complicated atmospheres. What we find is the following: (1) Although their numbers are small, it appears that the WRs in the SMC come from only the highest mass ($>70 \\cal M_\\odot$) stars. This is in accord with our expectations that at low metallicities only the most massive and luminous stars will have sufficient mass-loss to become WRs. (2) In the LMC, the early-type WN stars (WNEs) occur in clusters clusters whose turn-off masses range from 30$\\cal M_\\odot$ to 100 $\\cal M_\\odot$ or more. This suggests that possibly all stars with mass $>30 \\cal M_\\odot$ pass through an WNE stage at LMC metallicities. (3) The one WC star in the SMC is found in a cluster with a turn-off mass of 70$\\cal M_\\odot$,  the same as for the SMC WNs. In the LMC, the WCs are found in clusters with turn-off masses of 45$\\cal M_\\odot$ or higher, similar to what is found for the LMC WNs.  Thus we conclude that WC stars come from essentially the same mass range as do the WNs, and indeed are often found in the same clusters.  This has important implications for interpreting the relationship between metallicity and the WC/WN ratio found in Local Group galaxies, which we discuss. (3) The LBVs in our sample come from very high mass stars ($>85 \\cal M_\\odot$), similar to what is known for the Galactic LBV $\\eta$~Car, suggesting that only the most massive stars go through an LBV phase.  Recently, Ofpe/WN9 stars have been implicated as LBVs after one such star underwent an LBV-like outburst.  However, our study includes two Ofpe/WN9 stars, BE~381 and Br~18,  which we find in clusters with much lower turn-off masses ($25-35 \\cal M_\\odot$). We suggest that Ofpe/WN9 stars are unrelated to ``true\" LBVs: not all ``LBV-like outbursts\" may have the same cause. Similarly, the B[e] stars have sometimes been described as LBV-like.  Yet, the two stars in our sample appear to come from a large mass range ($>30-60 \\cal M_\\odot$). This is consistent with other studies suggesting that B[e] stars cover a large range in bolometric luminosities. (4) The bolometric corrections of early WN and WC stars are found to be extreme, with an average BC(WNE)=$-6.0$~mag, and an average BC(WC4)=$-5.5$~mag. These values are considerably more negative than those of even the hottest O-type stars.  However, similar values have been found for WNE stars by applying Hillier's ``standard model\" for WR atmospheres.  We find more modest BCs for the Ofpe/WN9 stars (BC=$-2$ to $-4$~mag), also consistent with recent analysis done with the standard model. Extension of these studies to the Galactic clusters will provide insight into how massive stars evolve at different metallicities. ",
        "introduction": "Conti (1976) first proposed that Wolf-Rayet (WR) stars might be a normal, late stage in the evolution of massive stars. In the modern version of the ``Conti scenario\" (Maeder \\& Conti 1994), strong stellar winds gradually strip off the H-rich outer layers of the most massive stars during the course of their main-sequence lifetimes. At first the H-burning CNO products He and N are revealed, and the star is called a WN-type WR star; this stage occurs either near the end of core-H burning or after core-He burning has begun, depending upon the luminosity of the star and the initial metallicity. Further mass-loss during the He-burning phases exposes the triple-$\\alpha$ products C and O, and results in a WC-type WR star. Since the fraction of mass that a star loses during its main-sequence evolution depends upon luminosity (mass), we would expect that at somewhat lower masses evolution proceeds only as far as the WN stage.  At still lower masses a star never loses sufficient mass to become a Wolf-Rayet at all, but spends its He-burning life as a red supergiant (RSG). Mass-loss rates also scale with metallicity as the stellar winds are driven by radiation pressure acting through highly ionized metal lines. Thus the mass-limits for becoming WN or WC stars should vary from galaxy to galaxy, and with location within a galaxy that has metallicity variations.  Studies of mixed-age populations in the galaxies of the Local Group have confirmed some of the predictions of the Conti scenario. For instance, the number ratio of WC and WN stars is a strong function of metallicity (Massey \\& Johnson 1998 and references therein), with proportionally more WC stars seen at higher metallicities, suggesting that the mass-limit for becoming WC stars is somewhat lower in these galaxies.  Similarly the relative number of WRs and RSGs is  correlated with metallicity, and there is a paucity of high luminosity RSGs at high metallicities (Massey 1998a), suggesting that these high luminosity stars have become WRs rather than RSGs.  However, fundamental questions remain concerning the evolution of massive stars:  \\noindent (1) What is the role of the luminous blue variables (LBVs)?  These stars are highly luminous objects that undergo photometric ``outbursts\" associated with increased mass-loss (Humphreys \\& Davidison 1994).  Are LBVs a short but important stage in the lives of {\\it all} high mass stars that occur at or near the end of core-H burning? Recent efforts have linked some of the LBVs to binaries, as Kenyon \\& Gallagher (1985) first suggested. The archetype LBV, $\\eta$ Car,  may be a binary with a highly eccentric orbit (Damineli, Conti, \\& Lopes 1997), but whether its outbursts have anything to do with the binary nature remains controversial (Davidson 1997), as does the orbit itself (Davidson et al.\\ 2000). Similarly, the WR star HD~5980 in the SMC underwent an ``LBV-like\" outburst (Barba et al 1995); this star is also believed to be a binary with an eccentric orbit, although the nature (and multiplicity?) of the companion(s) remains unclear (Koenigsberger et al.\\ 1998; Moffat 1999).  \\noindent The Ofpe/WN9 type WRs, and the high-luminosity B[e] stars have recently been implicated in the LBV phenomenon. The former have spectral properties intermediate between ``Of\" and ``WN\" (Bohannan \\& Walborn 1989). One of the prototypes of this class, R~127, underwent an LBV outburst in 1982 (Walborn 1982; Stahl et al.\\ 1983; see discussion in Bohannan 1997). Similarly some B[e] stars have been described as having LBV-like outbursts.  Var~C, a well-known LBV in M~33, has a spectrum indistinguishable from B[e] stars: compare Fig.~8a of Massey et al.\\ (1996) with Fig.~8 of Zickgraf et al. (1986). Do all B[e] stars undergo an LBV phase or not? Conti (1997) has provided an insightful review.  \\noindent (2) What is the evolutionary connection between WN and WC stars? We expect only the highest mass stars become WCs, while stars of a wider range in mass become WNs.  The changing proportion of WCs and WNs within the galaxies of the Local Group have been attributed to the expected dependence of these mass ranges on metallicity.  However, the relative time spent in the WN and WC stages may also change with metallicity, complicating the interpretation of such global measures drawn from mixed-age populations.  \\noindent (3) Is there any evolutionary significance to the excitation subtypes? Both WN and WC stars are subdivided into numerical classes, or more coarsely into ``early\" (WNE, WCE) or ``late\" (WNL, WCL) based upon whether higher or lower excitation ions dominate. Recent modeling by Crowther (2000) suggests that the distinction between WNL and WNE is not actually due to temperature differences but primarily metal abundance. Armandroff \\& Massey (1991) and Massey \\& Johnson (1998) have argued that this true for the WC excitation classes based upon the metallicity of the regions where these stars are found.  If we knew the progenitor masses of LBVs and the various kinds of WRs we would have our answers to the above. However, here recourse to stellar evolution models fails us. Stellar evolutionary models show that a star's path in the HRD during core-He burning is strongly dependent upon the amount of mass-loss that has preceded this stage. Thus the nature of the LBV phenomenon becomes very important in understanding where WRs come from, as the amount of mass ejected by LBVs is large, but given the episodic nature of LBVs, hard to include in the evolutionary models. In addition, the locations of WRs and LBVs in the H-R diagram are highly uncertain. LBVs have pronounced UV-excesses and ``pseudo-photospheres\" (Humphreys \\& Davidson 1994). For WR stars,  neither the effective temperatures nor bolometric corrections are established, as none of the standard assumptions of stellar atmospheres hold in the non-LTE, rapidly expanding, ``clumpy\" stellar winds where both the stellar continua and emission-lines arise (e.g., Conti 1988). While the WR subtypes represent some sort of excitation sequence in the stellar winds, the relationship, if any, to the effective temperature of the star remains unclear.  There has been recent success in modeling WR atmospheres, with convincing matches to the observed line profiles and stellar continua from the UV to the near-IR. These models have the potential for determining the bolometric luminosities and effective temperatures. The ``standard WR model\" (Hillier 1987, 1990) assumes a spherical geometry and homogeneity, and then iteratively solves the equations for statistical equilibrium and radiative equilibrium for an adopted velocity law, mass-loss rate, and chemical composition. (See also Hillier \\& Miller 1998, 1999.) Comparison with observations then permits tweaking of the parameters.  Although the solutions may not be unique, good agreement is often achieved with observations, and in a series of papers, Crowther and collaborators have offered the ``fundamental\" parameters (effective temperatures, luminosities, chemical abundances, mass-loss rates, etc.) of WN stars obtained with this model (Crowther, Hillier \\& Smith 1995a, 1995b;  Crowther, Smith, \\& Hillier 1995c; Crowther et al. 1995d; Crowther, Smith, \\& Willis 1995e; Crowther \\& Smith 1997; Bohannan \\& Crowther 1999).  Here we utilize a complementary, observational approach to the problem, one that can not only answer the question of the progenitor masses of LBVs and WRs, but also provide data on the BCs that can help constrain and evaluate the WR atmosphere models.  \\subsection{The Use of Cluster Turn-offs}  A time-honored method of understanding the nature of evolved stars is to determine the turn-off luminosities in clusters containing such objects (Johnson \\& Sandage 1955; Schwarzschild 1958). This was first applied by Sandage (1953) to determine the masses of RR~Lyrae stars in the globular clusters M~3 and M~92, with a result that was at variance with that given by theory (Sandage 1956). Similarly, the turn-off masses of intermediate-age open clusters were used by Anthony-Twarog (1982) to determine the progenitor masses of white dwarfs. However, it is one thing to apply this to clusters with ages of $10^{10}$ yr, as was done for the RR~Lyrae stars, or to clusters whose ages are $2\\times 10^{7}$---$7\\times 10^{8}$ yr, as was done for white dwarfs. Can we safely extend this to clusters whose ages are only of order 3--5$\\times 10^6$~yr in order to determine the progenitor masses of WRs and LBVs?  When stars form in a cluster or association, stars of intermediate mass appear to form over a significant time span---perhaps over several million years (Hillenbrand et al.\\ 1993; Massey \\& Hunter 1998). However, modern spectroscopic and photometric studies have shown that the massive stars tend to form in a highly coeval fashion.  For instance, in their study of the stellar content of NGC~6611, Hillenbrand et al.\\ (1993) found a {\\it maximum} age spread of 1~Myr for the massive stars, and noted that the data were consistent with {\\it no} discernible age spread. for all one could tell ``the highest-mass stars could have all been born on a particular Tuesday.\"  Similarly, the high mass stars in the R136 cluster have clearly formed over $\\Delta \\tau < 1$~Myr, given the large number of O3~V stars and the short duration that stars would have in this phase (Massey \\& Hunter 1998).  Such short time scales for star formation are consistent with recent studies by Elmegreen (1997, 2000a, 2000b), who argues that star formation takes place not over tens of crossing times but over one or two. For regions with large spatial extent (such as 100~pc diameter OB associations) star formation in the general region may occur over a prolonged time ($\\leq$10~Myr). However, large OB associations can contain subgroups that have formed independently (Blaauw 1964), and are small enough so that a high degree of coevality ($<1-2$~Myr) is expected.  The stars from such a subgroup need not be spatially coincident.  Rather, a star with a random motion of 10~km~s$^{-1}$ will have traveled 30~pc in just 3~Myr. Thus in an OB association we may find intermediate-mass stars which have formed from a number of subgroups over time, but  massive stars which may have formed from a single subgroup and hence are coeval---even though these massive stars may now be spread out throughout the OB association. Or, it may be that massive stars of different ages are present, in which case the ``turn-off mass\" will not be relevant to the evolved object.  We take an optimistic approach in our search for turn-off masses, but will insist that coevality be established empirically for the massive stars in the region in question.  For massive stars, the mass-luminosity relationship is much flatter than for solar-type stars ($L\\sim M^{2.4}$ for 30~$\\cal M_\\odot$ and $L\\sim M^{1.5}$ for 120~$\\cal M_\\odot$). As a result, the lifetimes of massive stars do not change as drastically with mass as one might expect.  A 120 $\\cal M_\\odot$ will have a main-sequence lifetime of 2.6~Myr, a 60 $\\cal M_\\odot$ still will have a main-sequence lifetime of 3.5~Myr, and a 25 $\\cal M_\\odot$ star will have a main-sequence lifetime of 6.4~Myr. (These numbers are based on the $z=0.02$ models of Schaller et al.\\ 1992.)  Thus it should be possible to use clusters and OB associations to pin down the ``minimum mass\" of various unevolved massive stars.  If the highest mass star still on the main-sequence is 60$\\cal M_\\odot$, and its associated stellar aggregate contains a WC-type WR star, then we might reasonably conclude that the progenitor mass of the WC star was at least 60 $\\cal M_\\odot$. Of course, if coevality does not hold, then this answer may be wrong---the WC star might have come from a 25 $\\cal M_\\odot$ that formed earlier.  But were that the case, it would have to have formed {\\it much} earlier---at least 3~Myr earlier, according to the lifetimes given above, and such an age spread should be readily apparent.  We can in principle also find the BCs from the cluster turn-offs. It is straightforward to determine the absolute visual magnitude of the WR, making some modest correction for the emission lines. Since massive stars evolve at nearly constant bolometric luminosity, we expect that the bolometric luminosity of the WR will be at least as great as the bolometric luminosity of the highest mass main-sequence object.  With modern stellar models we can improve on this by making first-order correction for modest luminosity evolution.  We are, of course, not the first to have trod on this ground. Schild \\& Maeder (1984) attempted to provide links between the different WR subtypes using this sort of analysis of Galactic clusters, concluding that stars with masses as low as $18 \\cal M_\\odot$ became WN stars, while WC stars came from stars of $35 \\cal M_\\odot$ and higher, and proposing various evolutionary relationships between the various subtypes. Humphreys, Nichols, \\& Massey (1985) also used data drawn from the literature on (mostly the same) Galactic clusters, and found a considerably higher minimum mass for becoming a WR star (30 $\\cal M_\\odot$), with no difference between the masses required to become a WN or a WC.  They were also the first to apply this method to determining the minimum bolometric corrections for WR stars, concluding that WNE stars have BCs $<-5.5$~mag, WNL stars have BCs $<-3.5$~mag, and WCs have BCs $<-5.0$~mag.  (These BCs are considerably more negative than had been commonly assumed.)  Smith, Meynet, \\& Mermilliod (1994) re-addressed the issue of BCs by analyzing the same data from the literature on what was also mostly the same clusters, finding BCs for WNs that were typically $-4$ mag (WNL) to $-6$ mag (WNE), and $-4.5$ for WCs, essentially unchanged from the Humphreys et al.\\ findings.  There were problems, however, with these earlier studies. The most significant one was the reliance upon (the same) literature data for the spectral types of the main-sequence stars in these clusters and associations. Over the past decade we have examined the stellar content of numerous clusters and OB associations in the Milky Way, and invariably discovered stars of high mass that had been previously missed either due to reddening or simple oversight (Massey, Johnson, \\& DeGioia-Eastwood 1995a). A related problem is that some of the literature spectral types were ``outdated\" for the O-type stars, particularly for stars of type O7 and earlier, which would lead to an incorrect assignment of bolometric corrections and hence luminosities and masses. In addition, our understanding of massive star evolution has improved to the point where we can do a considerably better job in assigning masses, and in particular understand the errors associated with this procedure (see, for example, Massey 1998b). Another problem was that the spectral information was sufficiently sparse that no test of coevality could be applied to the cluster.  In addition, poor photometry---often photographic---led to poor reddening corrections. And, finally, a significant limitation in these earlier studies was that all were restricted to the Milky Way. It would be most interesting to understand the origin of evolved massive stars as a function of metallicity; for this, extension to the Magellanic Clouds is a logical step.  We have attempted to rectify these problems by carrying out a modern analysis of OB associations containing WR and other evolved massive stars in galaxies of the Local Group, obtaining new spectroscopic and photometric data where warranted, and combining this with studies drawn form the recent literature. In this first paper we will determine the progenitor masses of WR and LBVs in 19 associations of the Magellanic Clouds. These two galaxies have abundances which are low compared to the solar neighborhood. In the next paper we will compare these to new results obtained for OB associations in our own Galaxy. In a third paper we will combine {\\it HST} photometric and spectroscopic data with large-aperture ground-based studies to extend this work to the more distant members of the Local Group as an addition check on metallicity effects.  Throughout this paper we will assume the true distance modulus of the SMC is 18.9, and that of LMC is 18.5 (Westerlund 1997; van den Bergh 2000). ",
        "conclusions": "Our photometric and spectroscopic investigation of 19 OB associations in the Magellanic Clouds has found that most of the massive stars have formed within a short time ($<$1~Myr) in about half of the regions in our sample. Their degree of coevality is similar to that found by Hillenbrand et al.\\ (1993) for NGC~6611, i.e., that the data are {\\it consistent} with all of the massive stars ``having been born on a particular Tuesday.\" In other regions, star-formation of the massive stars may have proceeded over a longer time, as suggested by the presence of evolved stars of 15-20$\\cal M_\\odot$ (suggesting ages of 10~Myr) along with unevolved stars of high mass (60 $\\cal M_\\odot$) with ages of only 2~Myr.  In some cases such apparent non-coevality may be due to chance line-of-sight coincidences within the Clouds, or to drift of lower mass stars into the space occupied by a truly coeval OB association, but in other cases, such as the $\\beta$ subcluster of LH~90, one is forced to conclude that star-formation itself was not very coeval, but proceeded over several million years.  The turn-off masses of the coeval associations have provided considerable insight into the evolution of massive stars.  We find that only the highest mass stars ($>70 \\cal M_\\odot$) become WRs in the SMC. The numbers are admittedly sparse, and an additional complication is the fact that most SMC WRs show the presence of absorption lines. Are these absorption lines indicative of a weak stellar wind (as evidenced by the weakness of the WR emission lines) or are these all due to binary companions? Conti et al.\\ (1989) discuss this without reaching any conclusions, and we note here that the issue of the binary frequency of the SMC WR stars requires further investigation.  Possibly a strong stellar wind due to very high luminosity {\\it and} binary-induced mass-loss is needed to become a WR star in the low metallicity of the SMC.  In the LMC the mass limit for becoming a WR star would appear to be a great deal lower, possibly 30$\\cal M_\\odot$.  Stars with a large range of initial masses (30-60 $\\cal M_\\odot$), and possibly {\\it all} massive stars with a mass above 30$\\cal M_\\odot$ go through a WNE stage in the LMC.  Most WR stars in the LMC are of early WN type; this is not true at the higher metallicity of the Milky Way, where WN3 and WN4 stars are relatively rare.  This is consistent with recent theoretical work of Crowther (2000), who finds that varying only the abundance in synthetic WN models (holding all other physical parameters consist) changes the spectral subtype, with WNEs characteristic of low abundances, and WNLs characteristic of higher abundances.  Thus, it may be the excitation classes are related neither to the masses nor to stellar temperatures.  The true LBVs occurs in clusters with very high turn-off masses ($\\approx 85\\cal M_\\odot$), both in the LMC and the SMC.  This is very similar to the turn-off mass in the Trumpler 14/16 complex with which the Galactic LBV $\\eta$~Car is associated (Massey \\& Johnson 1993).  This supports the standard picture, that LBVs are an important, if short-lived, phase in the evolution of the most massive stars, at least at the metallicities that characterize the Magellanic Clouds and the Milky Way.  We note with interest the important study by King, Gallagher, \\& Walterbos (2000), who find that some LBV stars in M~31 appear to be found in relative isolation, leading them to question whether these are all high mass stars, at least at the higher metallicity of M~31.  The Ofpe/WN9 stars, some of which go through some sort of outburst, cannot be ``true\" LBVs, if the nature of the latter is tied to extremely high bolometric luminosities.  We find that the Ofpe/WN9 stars have the {\\it lowest} masses of {\\it any} WRs, with the progenitors possibly as low as 25$\\cal M_\\odot$.    Similarly, the connection of the B[e] stars to LBVs seems tenuous on the basis of mass or bolometric luminosities.  We know that the relative number of WC and WN stars change drastically throughout the Local Group, in a manner well-correlated with metallicity (Massey \\& Johnson 1998). One obvious interpretation of this is that it is much harder to lose enough mass to become a WC star in a low-metallicity environment; i.e., only the most luminous and massive stars have sufficiently high mass-loss rates to achieve this.  And, similarly, the limit for WN stars should be higher in lower metallicity systems.  As long as the bar is somewhat lower for achieving WN status compared to WC status, then the IMF assures that the WC/WN ratio will change.  Thus our finding here that WCs and WNs come from similar mass ranges (although higher in the SMC than in the LMC), suggest that an alternative explanation is needed. Instead, it may be that it is the relative lifetimes of the WC and WN stages which are different at different masses; i.e., at very high masses the WC stage is shorter compared to the length of the WN stage than at lower masses.  Or, it could be that the metallicity itself affects the relative lifetimes of the WC and WN stages.  We note that we found luminous red supergiants (RSGs) cohabiting with both WNs and WCs in many OB associations in more distant galaxies of the Local Group (Massey \\& Johnson 1998; see for example their Figs.~14-16).  While we were unable to evaluate the degree of coevality of these associations, the statistics suggest that these stars have similar progenitor mass at a given metallicity, and that variations in the relative number of RSGs to WRs are due primarily to changes in the relative lifetimes due to the effect of metallicity on the mass-loss rates (Azzopardi, Lequeux, \\& Maeder 1988).  We conclude that the BCs of WNE stars are quite substantial, $-6$ mag.  This value is in very good accord with that determined from weak-lined WNE stars using the WR ``standard model\" of Hillier (1987, 1990) by Crowther et al.\\ (1995c). The earliest-type WN star known (of type WN2) is included in our sample, and our data suggest an even more striking BC ($<-7.5$~mag); a full analysis of Br~4 via the standard model would be of great interest. For the Ofpe/WN9 stars we find BCs of $-2$ to $-4$~mag, again in good agreement with the atmospheric analysis of several such stars by Crowther et al.\\ (1995a).  We find here that the BCs of WC4 stars are typically about $-5.5$~mag.  In the next paper, we will extend this study to the higher metallicities found in our own Milky Way galaxy."
    },
    "0002/astro-ph0002143_arXiv.txt": {
        "abstract": "We investigate the formation of binary stellar systems. We consider a model where a `seed' protobinary system forms, via fragmentation, within a collapsing molecular cloud core and evolves to its final mass by accreting material from an infalling gaseous envelope.  This accretion alters the mass ratio and orbit of the binary, and is largely responsible for forming the circumstellar and/or circumbinary discs.  Given this model for binary formation, we predict the properties of binary systems and how they depend on the initial conditions within the molecular cloud core.  We predict that there should be a continuous trend such that closer binaries are more likely to have equal mass components and are more likely to have circumbinary discs than wider systems.  Comparing our results to observations, we find that the observed mass-ratio distributions of binaries and the frequency of circumbinary discs as a function of separation are most easily reproduced if the progenitor molecular cloud cores have radial density profiles between uniform and $1/r$ (e.g. Gaussian) with near uniform-rotation.  This is in good agreement with the observed properties of pre-stellar cores.  Conversely, we find that the observed properties of binaries cannot be reproduced if the cloud cores are in solid-body rotation and have initial density profiles which are strongly centrally condensed. Finally, in agreement with the radial-velocity searches for extra-solar planets, we find that it is very difficult to form a brown dwarf companion to a solar-type star with a separation $\\simless 10$ AU, but that the frequency of brown dwarf companions should increase with larger separations or lower mass primaries. ",
        "introduction": "\\label{introduction}  The favoured mechanism for producing most binary stellar systems is the fragmentation of a molecular cloud core during its gravitational collapse.  Fragmentation can be divided into two main classes: direct fragmentation (e.g.~Boss \\& Bodenheimer 1979; Boss 1986; Bonnell et al.~1991, 1992; Bonnell \\& Bastien 1992; Nelson \\& Papaloizou 1993; Burkert \\& Bodenheimer 1993; Bate \\& Burkert 1997), and rotational fragmentation (e.g.~Norman \\& Wilson 1978; Bonnell 1994; Bonnell \\& Bate 1994a, 1994b; Burkert \\& Bodenheimer 1996; Burkert, Bate, Bodenheimer 1997). Direct fragmentation depends critically on the initial density structure within the molecular cloud core (e.g.~non-spherical shape or density perturbations), whereas rotational fragmentation is relatively independent of the initial density structure of the cloud because the fragmentation occurs due to nonaxisymmetric instabilities in a massive rotationally-supported disc or ring.  The main conclusion, from $\\approx 20$ years of fragmentation studies, is that it appears to be possible to form binaries with similar properties to those that are observed. However, it has not been possible to use these calculations to predict the fundamental properties of stellar systems such as the fraction of stellar systems which are binary or the properties of binary systems (e.g.~the distributions of mass ratios, separations, and eccentricities and the properties of discs in pre-main-sequence systems).  There are two primary reasons for this lack of predictive power. First, the results of fragmentation calculations depend sensitively on the initial conditions, which are poorly constrained. The second problem is that of accretion.  In fragmentation calculations, the binary or multiple protostellar systems that are formed initially contain only a small fraction of the total mass of the original cloud (e.g.~Boss 1986; Bonnell \\& Bate 1994b) with the magnitude of this fraction decreasing with the binary's initial separation (see Section \\ref{bmass_vs_sepsec}). To obtain the final parameters of a stellar system, a calculation must be followed until all of the original cloud material has been accumulated by one of the protostars or their discs. Unfortunately, due to the enormous range in densities and dynamical time-scales in such a calculation, this is very difficult. Thus far, only one calculation has followed the three-dimensional collapse of a molecular cloud core until $>90$\\% of the initial cloud was contained in the protostars or circumstellar/circumbinary discs (Bate, Bonnell \\& Price 1995). Because such calculations are so difficult to perform, it is impossible to perform the number of calculations that would be required to predict the statistical properties of binary stellar systems -- even if we knew the distribution of initial conditions.  On the other hand, if we can overcome this second difficulty, we can use observations of binary systems to better constrain the initial conditions for star formation.  Bate \\& Bonnell \\shortcite{BatBon97} quantified how the properties of a binary system are affected by the accretion of a small amount of gas from an infalling gaseous envelope. They found that the effects depend primarily on the specific angular momentum of the gas and the binary's mass ratio (see also Artymowicz 1983; Bate 1997).  Generally, accretion of gas with low specific angular momentum decreases the mass ratio and separation of the binary, while accretion of gas with high specific angular momentum increases the separation, drives the mass ratio toward unity, and can form a circumbinary disc. From these results, they predicted that closer binaries should have mass ratios that are biased toward equal masses compared to wider systems.  In this paper, we use the results of Bate \\& Bonnell \\shortcite{BatBon97} to develop a protobinary evolution code that enables us to follow the evolution of a protobinary system as it accretes from its initial to its final mass, but does so in far less time than would be required for a full hydrodynamic calculation. Using this code, we consider the following model for the formation of binary stellar systems.  We assume that a `seed' binary system is formed at the centre of a collapsing molecular cloud core, presumably via some sort of fragmentation.  The protobinary system initial consists of only a small fraction of the total mass of the core. Subsequently, it accretes the remainder of the initial cloud (which is falling on to the binary) and its properties evolve due to the accretion.  We consider the formation process to be complete when all of the original cloud's material is contained either in one of the two stars or their surrounding discs. Our goal is to obtain predictions about the properties of binary stars that can be tested observationally, and to determine how these properties depend on the initial conditions (e.g.~the density and angular momentum profiles) in the progenitor molecular cloud cores so that the initial conditions can be better constrained.  In Section \\ref{commethod}, we describe the methods used to follow the evolution of accreting protobinary systems, and we present the results of various test calculations in Section \\ref{comparison}. In Sections \\ref{evolution} and \\ref{relax}, we present results from calculations with a range of initial conditions which follow the evolution of accreting protobinary systems. From these results, in Section \\ref{predictions}, we make predictions regarding the properties of binary systems and compare them with the latest observations. These predictions are briefly summarised in Section \\ref{conclusions}.  Those readers more interested in our predictions of the properties of binary stars, rather than the method by which these predictions have been obtained, may care to move directly to Section \\ref{evolution}. ",
        "conclusions": "\\label{conclusions}  We have considered a model for the formation of binary stellar systems which has been inspired by the results obtained from $\\approx 20$ years of study of the fragmentation collapsing molecular cloud cores. In the model, a `seed' protobinary system forms, presumably via fragmentation, within a collapsing molecular cloud core and evolves to its final mass by accreting material from an infalling gaseous envelope.  We developed and tested a method which can rapidly follow the evolution of the mass ratios, separations and circumbinary disc properties of such binaries as they accrete to their final masses. Using this protobinary evolution code, we predict the properties of binary stars and how they depend on the pre-collapse conditions in their progenitor molecular cloud cores.  These predictions and their comparison with current observations are discussed in detail in Section \\ref{predictions}.  Briefly, we conclude that, if most binary stars form via the above model, binary systems with smaller separations or greater total masses should have mass ratios which are biased toward equal masses when compared to binaries with wider orbits or lower total masses.  Similarly, the frequency of circumbinary discs should be greater for pre-main-sequence binaries with closer orbits or greater total masses. These conclusions can be understood because: binaries which are closer or have a greater final mass should accrete more gas relative to their initial masses than wider or lower-mass binaries; the specific angular momentum of the infalling gas relative to that of the binary is expected to increase as the accretion proceeds; and the accretion of gas with high specific angular momentum tends to equalise the mass ratio and forms a circumbinary disc. We also demonstrate that in a young binary which is accreting from an infalling gaseous envelope, the primary will generally have a circumstellar disc which is more massive or similar in mass to that of the secondary.  All of these conclusions are in good agreement with the latest observations.  By making rough quantitative predictions of the properties of binary stars, we find that the observed properties of binary stars are most easily reproduced if the pre-collapse molecular cloud cores from which binaries form have radial density profiles between uniform and $1/r$ (e.g.~Gaussian) with near uniform rotation.  This is in excellent agreement with the observed properties of pre-stellar cores (Ward-Thompson et al.~1994; Andr\\'e, Ward-Thompson, Motte 1996; Ward-Thompson, Motte, \\& Andr\\'e 1999). Conversely, the observed properties of binaries cannot be reproduced if the cores are in solid-body rotation and have initial density profiles which are strongly centrally condensed (between $1/r$ and $1/r^2$), and the singular isothermal sphere ($\\rho \\propto 1/r^2$) cannot fit the observations even with strong differential rotation."
    },
    "0002/astro-ph0002469_arXiv.txt": {
        "abstract": "We report results from a Hubble Space Telescope (HST) and Near-Infrared Camera and Multiobject Spectrometer (NICMOS) program to study the distribution of hot neutral (molecular hydrogen) and ionized circumstellar material in the young planetary nebulae \\n7027. HST/NICMOS provided very high spatial resolution imaging in line and continuum emission, and the stability and large dynamic range needed for investigating detailed structures in the circumstellar material. We present dramatic new images of \\n7027 that have led to a new understanding of the structure in this important planetary nebula. The central star is clearly revealed, providing near-infrared fluxes that are used to directly determine the stellar temperature very accurately (T$_{\\star}$ = 198,000 K). It is found that the photodissociation layer as revealed by near--infrared molecular hydrogen emission is very thin ($\\Delta$R $\\sim$ 6$\\times$10$^{15}$ cm), and is biconical in shape. The interface region is structured and filamentary, suggesting the existence of hydrodynamic instabilities. We discuss evidence for the presence of one or more highly collimated, off-axis jets that might be present in \\n7027. The NICMOS data are combined with earlier Hubble Space Telescope data to provide a complete picture of NGC 7027 using the highest spatial resolution data to date. The evolutionary future of \\n7027 is discussed. ",
        "introduction": "Planetary nebulae have long been thought of as ionized remnants of circumstellar material ejected by stars on the asymptotic giant branch (AGB). That somewhat limited view has changed dramatically, as the neutral component of the remnant circumstellar envelope has been shown to remain observable very late into the lifetime of many planetary nebulae (PNe; \\eg \\cite{din95}; \\cite{hug96}; \\cite{hor99}; and references therein). It is now known that many PNe are characterized by emission from material that is contained in rapidly evolving photodissocation regions (PDRs; \\cite{th85}; \\cite{sd89}; \\cite{lat93a}; \\cite{nat98}). As the central star and nebula evolve, UV emission from the star increases rapidly. A photodissociation front moves through the gas and slowly turns the mostly molecular nebula to predominantly atomic. At this point, the nebula radiates mainly in the infrared. When the central star becomes hotter than $T_* \\approx 30,000$ K, the circumstellar gas quickly becomes ionized and the planetary nebula shines brightly in visible light. The chemical properties of gas in PNe and proto-PNe (PPN) are like those of interstellar PDRs (\\cite{lat92}; \\cite{lat93a}; \\cite{tie93}; \\cite{hor94}), which are well modeled (see \\cite{ht97}).  The time over which molecular material can be present in PNe is typically predicted to be a rather brief period in the lifetime of PNe (\\eg \\cite{tie93}; \\cite{nat98}). But, such timescales predicted from current chemical models are limited by observational data and lack of a clear understanding of the morphology and density structure of these objects.  \\n7027 is perhaps the best studied planetary nebula, in part because of its proximity (about 900 pc; e.g. \\cite{mas89}) and high surface brightness at all wavelengths. It has a very rich atomic and molecular spectrum, making it ripe for study at all wavelengths, and it is rich in physical and chemical information (see, \\eg \\cite{gra93a}; \\cite{cox97}, and references therein). Because of its compact size as viewed from Earth ($\\approx 15$\\arcsec), the morphology, particularly that for the \\h2\\ emission, of this object has been suggested, but has not been clearly revealed (see, \\eg \\cite{gra93a}; \\cite{gra93b}; \\cite{kas94}; \\cite{lat95}). The elliptical ionized core lies at the center of an extended molecular envelope (\\cite{jhb91}). At the interface between the cold molecular envelope and the hot ionized region is an apparent ``quadrupolar'' region of strong near--infrared (IR) rotational-vibrational molecular hydrogen emission (\\cite{gra93a}) that has been shown to be excited by absorption of UV photons in the photodissociation region (\\cite{hor99}).  A detailed understanding of the morphology of planetary nebulae is important. It has been demonstrated that there is a strong correlation between the presence of molecular emission from PNe and the observed morphology of the PNe, such that objects that contain large amounts of molecular material are bipolar or butterfly nebulae (\\cite{hor99}; \\cite{kas96}; \\cite{hug96}; \\cite{hug89}; \\cite{zuc88}). Progenitor mass also correlates with morphological type, such that higher mass stars appear to produce bipolar or butterfly nebulae (see \\cite{cor94}). This correlation suggests that the higher mass AGB progenitors, which probably have the highest mass loss rates, produce dense, long lived molecular envelopes (\\cite{hor99}).  The evolution of PDR and PNe, and the correlation between morphology with molecular content are not fully understood. Because NGC 7027 is in a rapid and key moment in evolution (the transition from neutral, predominantly molecular envelope to an ionized one), it is worthy of serious study into its chemical and physical properties and its detailed morphology.  The Hubble Space Telescope and NICMOS provided the high spatial resolution, dynamic range, and stable background needed to examine the near-IR nature of NGC 7027. In this paper, we present narrowband images in the \\s1 ($\\lambda = 2.121$ \\micron) line of molecular hydrogen and other filters that trace the ionized core and the nearby neutral region. These images show with unprecedented clarity the true structure of NGC 7027. In Sections 2 and 3, we discuss the data, the apparent morphology, and the presence of a jet (or jets) in the object. In \\S 4 we present a 3-dimensional model for the overall structure, including the photodissociation region. Section 5 presents a discussion of the excitation of the nebula and the central star properties. We then consider our results in light of previous work, and discuss the evolution of the PDR in NGC 7027. ",
        "conclusions": "\\subsection{Previous Kinematic Results}  In a detailed multiwavelength study, Graham \\etal (1993a) found that uniform, radial expansion of a prolate spheroid, tilted to the line of sight is consistent with the millimeter-wave (CO) position-velocity maps. However the CO P-V diagram differs from the H$_2$ P-V diagram presented here. This is not unexpected since the observed emission from the two molecules is produced under very different excitation conditions. Therefore the two molecules trace different spatial/kinematical structures. Observations have been made with the ``BEAR'' instrument (a 256$\\times$256 pixel HgCdTe camera coupled to a Fourier transform spectrometer, with a spectral resolution \\about\\ 52.2 \\kms) on the Canada-France-Hawaii Telescope (\\cite{cox97}). Cox \\etal found that the velocity field appeared to them as an equatorial ``torus'' suggests that north pole is tilted away from observer. This result is inconsistent with our interpretation of the HST images, our 3-D modeling of those images, and the higher spectral resolution CSHELL data. The high spatial resolution and sensitivity of the HST images shows that the front part of the top rim lies in front of the the ionized nebula (not behind), suggesting that the top (northern side) is tilted $toward$ the observer. We cannot resolve this discrepancy with the data available. We suggest that the spectral resolution available to Cox \\etal (1997) was not sufficent to fully deconvolve the subtle, but complex velocity structure that must be present (\\cite{kel99}). In addition, the model schematic given by Jaminet \\etal (1991) shows the same orientation as Cox \\etal (north pole away) for the fast molecular wind traced by CO $J = 3 \\to 2$ line emission, which presumably has a similar flow direction as the ionized gas. It is entirely possible that the flow along the back of the cone (or bubble) has been confused with a flow along the polar axis, which in our model would not be the dominant emitting region. Bains \\etal (1997) found from their high spectral resolution radio and optical observations, an orientation for NGC 7027 that is consistent with ours: i.e. that the NE lobe is blue-shifted, SW lobe is red-shifted.  \\subsection{Evolution of the PDR in \\n7027}  A full modeling of PDR evolution is beyond the scope of this paper. We can, however, make an assessment of such PDR evolution based on available models. The morphological evolution of the PDR will depend most strongly on the distribution of material in the circumstellar envelope. For \\n7027\\ the post-AGB circumstellar envelope must have been axially symmetric with a gradual decrease in density from the equatorial region to the poles. A dense equatorial disk cannot be ruled out, but there is no evidence in these data, or other molecular data (\\cite{jhb91}; \\cite{jam91}; \\cite{gra93a}), that requires or even suggests one. For that reason, we consider such a structure as unlikely to be present. The density distribution appears typical for other post-AGB objects with an equatorial to polar density contrast of $\\sim 2 - 10$, such that it is not as great as the extreme bipolar proto-PNe (see \\S \\ref{h2e}). In much less than the cooling timescale of the central star, the PDR will be established quickly in all directions, with the fastest evolution in the polar regions. The result will most likely be for \\n7027\\ to become a butterfly-type nebula similar in gross properties to (e.g.)  NGC 2346. Wind interactions for this type of morphological evolution are not required.  The timescale for PDR evolution is highly uncertain (other timescales are discussed in \\S \\ref{cs}), but reasonable estimates can, and have been made (see, \\eg \\cite{tie93}). It is straightforward to show that for a constant mass loss rate and parameters typical of high mass loss rate carbon stars, the timescale for a PDR to completely move through the circumstellar envelope can be estimated by integrating the equations of molecular formation and destruction with respect to radius and time, as was done by \\cite{tie93}. For a constant wind velocity $V_w = 20$ \\kms, a constant mass loss rate $\\dot M = 10^{-4}$ \\msol\\ yr$^{-1}$, and a stellar luminosity $L_* = 6000$ \\lsol, it can be shown that the dissociation front travels a distance $r_i \\approx 3\\times 10^{16}$ cm in $t_i \\approx 65\\ {\\rm years}$ with a PDR speed of $V_i \\approx 263\\ {\\rm km\\ s}^{-1}$ (see \\cite{tie93} for a more detailed discussion; see also \\cite{nat98}). Tielens suggested clumping as how the evolution is slowed in real nebulae. Although widespread clumping and structure is seen in these data, it does not appear at the size scales required, or marginally so at best ($R \\gtrsim 10^{16}$ cm and $A_V \\gtrsim 4$ mag). Even if clumping is important in localized regions, it is more likely that we now know several of the assumptions used to derive this evolutionary timescale are not valid, and if properly accounted for will tend to increase the timescale. The mass loss rate is not constant. The ``ring'' structure clearly visible in AFGL 2688 (\\cite{lat93}; \\cite{sah98}) indicates a time varying mass loss rate on a timescale of hundreds of years. The same type of structures are seen in other objects observed (\\cite{kwok98}), including \\n7027 (see, \\eg \\cite{bond97}). In addition, the distribution of material is not spherical. The degree to which the mass loss rate is varying is not known. Nor is how a non-spherical distribution of attenuating and shielding material will alter the evolutionary timescales (see, however, \\cite{nat98}). Additional work must be done to characterize these properties and how they relate to evolution. From the above expression and an assessment of the impact of uncertain parameters, it seems likely that the timescale for \\n7027\\ to evolve away from its current state will be only a few hundred to a few thousand years (see also \\cite{nat98}). Wind interactions will complicate the picture by adding hydrodynamic effects and changes to the density structure. However, evolution to the current epoch appears to be described fully by the interaction of FUV and ionizing photons with an AGB wind that has (or had) a somewhat enhanced mass loss rate in the equatorial plane. While there is evidence for recent jet interactions with the outflow, the jets have not dominated the morphological shaping. It is worth monitoring \\n7027\\ at very high spatial resolution for morphological changes caused by jets and UV photons over the next several decades."
    },
    "0002/astro-ph0002375_arXiv.txt": {
        "abstract": "ROSAT HRI observations show complicated substructure in the X--ray surface brightness within $\\sim$5 arcminutes around NGC 1275 -- the dominant galaxy of the Perseus cluster. The typical amplitude of the variations is of the order of 30\\% of the azimuthally averaged surface brightness at a given distance from NGC 1275. We argue that this substructure could be related to the activity of NGC 1275 in the past. Bubbles of relativistic plasma, inflated by jets, be forced to rise by buoyancy forces, mix with the ambient intracluster medium (ICM), and then spread. Overall evolution of the bubble may resemble the evolution of a hot bubble during a powerful atmospheric explosion. From a comparison of the time scale of the bubble inflation to the rise time of the bubbles and from the observed size of the radio lobes which displace the thermal gas, the energy release in the relativistic plasma by the active nucleus of NGC 1275 can be inferred. Approximate modeling implies a nuclear power output of the order of $10^{45}$ erg s$^{-1}$ averaged over the last $\\sim 3~10^7$ years. This is comparable with the energy radiated in X-rays during the same epoch.  Detailed measurements of the morphology of the X--ray structure, the temperature and abundance distributions with Chandra and XMM may test this hypothesis. ",
        "introduction": "The Perseus cluster of galaxies (Abell 426) is one of the best studied clusters, due to its proximity ($z=0.018$, $1'$ corresponds to $\\sim$ 30 kpc for $H_0=50~km~s^{-1}~Mpc^{-1}$) and brightness. Detailed X--ray images were obtained with the Einstein IPC (Branduardi--Raymont et al. 1981) and HRI (Fabian et al. 1981) and the ROSAT PSPC (Schwarz et al. 1992, Ettori, Fabian, White 1999) and HRI (B\\\"{o}hringer et al. 1993; see also Heinz et al. 1998). The cluster has a prominent X--ray surface brightness peak at its center along with cool gas, which is usually interpreted as due to the pressure induced flow of gas releasing its thermal energy via radiation. The cooling flow is centered on the active galaxy NGC1275, containing a strong core-dominated radio source (Per A, 3C 84) surrounded by a lower surface brightness halo (e.g. Pedlar et al. 1990, Sijbring 1993). Analysis of the ROSAT HRI observations of the central arcminute has shown that the X-ray emitting gas is displaced by the bright radio emitting regions (B\\\"{o}hringer et al. 1993), suggesting that the cosmic ray pressure is at least comparable to that of the hot intracluster gas. Many other studies explored correlations of X-ray, radio, optical, and ultraviolet emission (see e.g. McNamara, O'Connell \\& Sarazin, 1996 and references therein). In this contribution, we discuss asymmetric structure in the X--ray surface brightness within $\\sim$ 5 arcminutes of NGC 1275 and suggest that buoyant bubbles of relativistic plasma may be important in defining the properties of this structure. ",
        "conclusions": "The X--ray surface brightness around NGC 1275 (dominant galaxy of the Perseus cluster) is perturbed at various spatial scales. We suggest that on arcminute scales, the disturbance is caused by bubbles of relativistic plasma, inflated by jets during the past $\\sim 10^8$ years. Overall evolution of the buoyant bubble will resemble the evolution of a hot bubble during a powerful atmospheric explosion. Colder gas from the central region of the cooling flow may be uplifted by the rising bubbles and (in the case of continuous jet activity) may make several cycles (from the center to the outer regions and back) on time scales comparable to the cooling time of the gas in the cooling flow.  A very important result that can be inferred from this model is the total power output of the nuclear energy source in NGC 1275 in the form of relativistic plasma. This energy release averaged over a time scale of about $3~10^7$ to $10^8$ years is estimated as a function of the inflation time of the central radio lobes, the rise time of the inflated bubbles due to buoyancy forces, and the actual size of the central bubbles. A geometrically simplified model yields a power output on the order of $10^{45}$ erg s$^{-1}$. This is comparable with the energy lost at the same time by thermal X-ray radiation from the entire central cooling  flow region. This raises the question, where does all this energy go, especially if the energy release is persistent over a longer epoch during which the relativistic electrons can lose their energy by radiation, but  the energy in protons and in the magnetic field is mostly conserved. The complicated X-ray morphology discussed in this paper may indicate long lasting nuclear activity, if we interpret the peculiar structure in the X-ray surface brightness as remnants of decaying  radio lobe bubbles.  Detailed measurements of the morphology of the X--ray structure and the temperature and abundance distribution with Chandra and XMM may test this hypothesis. The gas uplifted from the central region is expected to be cooler than the ambient gas and to have an abundance of heavy elements typical of the innermost region. If cosmic rays are mixed with the thermal gas, then the pressure, as derived from X--ray observations, may be lower than the pressure of the ambient gas."
    },
    "0002/astro-ph0002139_arXiv.txt": {
        "abstract": "s{ Type Ia supernovae (SNIa) have been  used as approximate standard candles to measure cosmological parameters such as the Hubble constant and the deceleration parameter. These measurements rely on empirical correlations between peak luminosities and other features that can be observed in the supernovae spectra and their light curves. Such correlations deserve further study since they have been established  from small samples of nearby SNIa. Two years ago, the EROS2 collaboration launched an automated search for supernovae with the 1m Marly telescope operating at La Silla. In all, 57 SNe have been  discovered in this EROS2 search and spectra have been  obtained for 26 of them. We found that 75\\% were of type Ia and 25\\% of type II. Using this sample, a preliminary SN explosion rate has been obtained. Our most recent observation campaign took place in February and March 99. It was performed  in the framework of a large consortium  led by the {\\em Supernova Cosmology Project}. The aim of this intensive campaign was to provide an independent set of high quality light curves and spectra to study systematic effects in the measurement of cosmological parameters. We will briefly describe our search procedure and present the status of our ongoing analysis. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002249_arXiv.txt": {
        "abstract": "I review the basic theory of the cosmic microwave background (CMB) anisotropies in adiabatic cold dark matter (CDM) cosmologies. The latest observational results on the CMB power spectrum are consistent with the simplest inflationary models and indicate that the Universe is close to spatially flat with a nearly scale invariant fluctuation spectrum. We are also beginning to see interesting constraints on the density of CDM, with a best fit value of $\\omega_c \\equiv \\Omega_c h^2 \\sim 0.1$. The CMB constraints, when combined with observations of distant Type Ia supernovae, are converging on a $\\Lambda$-dominated Universe with $\\Omega_m \\approx 0.3$ and $\\Omega_\\Lambda \\approx 0.7$.\\footnote{To appear in Proceedings of NATO ASI: Structure formation in the Universe, eds. N. Turok, R. Crittenden.} ",
        "introduction": "The discovery of temperature anisotropies in the CMB by the COBE team (Smoot \\etal 1992) heralded a new era in cosmology.  For the first time COBE provided a clear detection of the primordial fluctuations responsible for the formation of structure in the Universe at a time when they were still in the linear growth regime. Since then, a large number of ground based and balloon borne experiments have been performed which have succeeded in defining the shape of the power spectrum of temperature anisotropies $C_\\ell$\\footnote{The power spectrum is defined as $C_\\ell = \\langle \\vert a_{\\ell m}\\vert^2 \\rangle$, where the $a_{\\ell m}$ are determined from a spherical harmonic expansion of the temperature anisotropies on the sky, $\\Delta T/T = \\sum a_{\\ell m} Y_{\\ell m}(\\theta, \\phi)$.}  up to multipoles of $\\ell \\sim 300$ clearly defining the first acoustic peak in the spectrum. Figure 1 shows a compilation of band power anisotropy measurements \\begin{equation} {\\Delta T_\\ell \\over T} = \\sqrt{ {1 \\over 2 \\pi} \\ell (\\ell + 1) C_\\ell} \\end{equation} that is almost up to date at the time of writing. The horizontal error bars show the multipole range probed by each experiment. The recent results from the VIPER experiment (Peterson \\etal 1999) and the Boomerang test flight (Mauskopf \\etal 1999) are not plotted because the exact window functions are not yet publically available. Neither are the published results from the Python V experiment (Coble \\etal 1999) which seem to be  discrepant with the other experiments particularly in the multipole range $\\ell \\simlt 100$. The points plotted in figure 1 are generally consistent with each other and provide strong evidence for a peak in the power spectrum at $\\ell \\sim 200$.  \\begin{figure} \\vspace{9cm} \\special{psfile=pgconf1.ps hscale=65 vscale=65 angle=-90 hoffset= -20 voffset=340} \\caption{Current constraints on the power spectrum of CMB temperature anisotropies. The error bars in the vertical direction show $1\\sigma$ errors in the band power estimates and the error bars in the horizontal direction indicate the width of the band. The solid line shows the best fit adiabatic CDM model with parameters $\\omega_b = 0.019$, $\\omega_c = 0.10$, $n_s = 1.08$, $Q_{10} = 0.98$, $\\Omega_m=0.225$, $\\Omega_\\Lambda = 0.775$.} \\end{figure}  In this introductory article, I will review briefly the theory of CMB anisotropies in adiabatic models  of structure formation and then discuss the implications of Figure 1 for  values of cosmological parameters. The literature on the CMB anisotropies has grown enormously over the last few years and it is impossible to do the subject justice in a short article. General reviews of the CMB anisotropies are given by Bond (1996) and Kamionkowski and Kosowsky (1999). A recent review on constraining cosmological parameters from the CMB is given by Rocha (1999). ",
        "conclusions": ""
    },
    "0002/astro-ph0002280_arXiv.txt": {
        "abstract": "We exploited the large areal  coverage offered by the Digitized Palomar Observatory Sky Survey to analyze  the outermost  regions of  the  galactic globular cluster   \\object{M~92} (\\object{NGC~6341}).  Two independent photometric reduction   programs (SKICAT  and  DAOPHOT)  were  used to construct  a  color-magnitude  diagram and a surface   density profile for this cluster, based on  J- and F-band DPOSS  plates. A strong similarity  has been found in  the performance of the  two  programs in  the low--crowded outermost cluster regions.  After removing the background contribution, we  obtained the cluster  outer surface density profile down    to a  surface brightness  magnitude    of $\\mu_{\\mathrm V} \\sim$  31 mag arcsec$^{-2}$ and matched it with the inner profile  of Trager et al. (\\cite{Tra95}). The profile shapes match very well:  since our data are uncalibrated, the  shift in magnitudes between  the profiles  has been also used to calibrate our profile. The analysis shows that the cluster has an extra tidal halo  extending out to $\\sim 30\\arcmin$  from the  cluster center at a $3~\\sigma$ level over the background  noise. This halo is revealed to be almost circular. ",
        "introduction": "The tidal radii  of  globular clusters  (GCs)  are important tools for understanding the complex interactions of GCs with the Galaxy.  In fact, they have traditionally been  used to study the mass distribution of the galactic halo (Innanen et al. \\cite{Inn83}), or to deduce GCs orbital parameters (Freeman \\& Norris \\cite{Fre81}; Djorgovski et al. \\cite{DJ96}).  Tidal  radii have usually been {\\it   estimated} (only in   few  cases {\\it  directly} measured),  by fitting  King models to    cluster density profiles   rarely measured from  the inner regions out to  the tidal radius,  because of the nature of the photographic material, that  prevented any measure in the cluster center, and the small  format of the first digital cameras. Only  in the last  few years, the advent of deep digitized sky  surveys and wide  field digital detectors has allowed us to  deal with   the overwhelming   problem  of contamination from field stars and  to probe the outer region of GCs directly (Grillmair et al. \\cite{Gri95}, hereafter G95; Zaggia et al. \\cite{Zag95}; Zaggia et al. \\cite{Zag97}; Lehman \\& Scholz \\cite{LS97}).   The study of tidal  tails in galactic satellites is gaining interest  for many applications related to the derivation    of  the galactic  structure and   potential,  the formation and evolution of the galactic halo, as well as the dynamical evolution  of the clusters  themselves.   Recent determinations of proper motion for some  globular clusters  with HIPPARCOS have  made it possible to  estimate   the orbital parameters  of  a  good  number of them (Dinescu et al. \\cite{Din99}).  This  helps to clarify  the nature and  structure   of tidal  extensions  in GCs.  In  principle, available  tools  to enhance cluster star   counts against field stars rely   on the color-magnitude  diagram   (CMD), proper motions, radial velocities, or a combination of the three techniques.  The application of these techniques to  GCs have led   to the discovery that tidal  or extra-tidal  material  is a common  feature: Grillmair (\\cite{Gri98}), for instance,  reported the discovery   of  tidal tails in   16 out  of 21 globular clusters.  Interestingly, signature  of the presence of tidal tails in GCs has also been found in four GC's  in M31  (Grillmair \\cite{Gri96}).  For galactic clusters, the discovery was made by using a selection in the  CMD of  cluster stars   on catalogs  extracted from digitized photographic datasets.  The CMD selection  technique is  an economical and powerful method to detect  GC  tails, since it significantly decreases the number of background and/or foreground objects.  \\begin{figure*}[t] \\resizebox{10cm}{!}{\\includegraphics{testa_fig1.ps}} % \\parbox[b]{55mm}{ \\caption{A  comparison  of objects   detected in the   inner parts  of \\object{M~92} by SKICAT (filled  triangles) and DAOPHOT (dots) for the two  different plates: $J$  ({\\it upper  panel})  and $F$ ({\\it  lower panel}).   The inner  circle (continuous) marks  the circular aperture where plate detections cannot be used.    The outer circle  (short-dash) marks  the annular region where  crowding correction  is important. In both panels  North is up and  East is to  the right.  The two diagonal bands  in the lower  panel indicate the  satellite tracks where SKICAT detects no objects.} \\label{fig1}} \\end{figure*}   In order  to test the  feasibility of a  survey of most GCs present in the Northern hemisphere, we applied  the CMD technique to the galactic globular  cluster  \\object{M~92} (NGC~6341), with the aim of measuring  the tidal radius  and searching for  the possible presence  of extra-tidal  material.   We used   plates  from the  Digitized  Second Palomar Sky Survey (hereafter DPOSS), in the framework of the CRoNaRio (Caltech-Roma-Napoli-Rio de Janeiro) collaboration (Djorgovski et al. \\cite{DJ97}, Andreon et al. \\cite{And97}, Djorgovski et al. \\cite{DJ99}). A previous account on this work was given in  Zaggia et al. (\\cite{Zag98}).  This is the first of a series of  papers dedicated to the subject --an  ideal application for this kind of all-sky surveys.  \\begin{figure*}[t] \\resizebox{10cm}{!}{\\includegraphics{testa_fig2.ps}} \\hfill \\parbox[b]{55mm}{ \\caption{Comparison between SKICAT   $M_{Core}$ and DAOPHOT   aperture magnitude in the filters $F$ ({\\it upper  panel}) and $J$ ({\\it lower panel}).The histograms of the distributions, expressed in percentage of the total, are reported at the right edge of the plots.} \\label{fig2}} \\end{figure*} ",
        "conclusions": "We investigated the presence and significance of  a tidal extension of the brightness profile of \\object{M~92}. The main results of our study are: \\begin{enumerate}  \\item The presence of an extra-tidal  profile extending out to $\\sim~0.5^{\\circ}$ from the cluster center, at a significance level of $3~\\sigma$ out to  $r~\\sim~2000~\\arcsec$.  We found no strong evidence for preferential  direction of  elongation  of the profile.  This  may imply that we are detecting the extra-tidal halo of evaporating stars, which will later form a tidal stream. Moreover, the tidal tail might be compressed along the line of sight --see, for instance, Fig.~18 of G95. In fact, G95 point out that tidal tails extend over enormous distances ahead and behind the cluster orbit,  and the volume density is subject to the open-orbit analogous of Kepler's third law:  near apogalacticon, stars  in the tidal  tail undergo  differential  slowing-down, so that  the  tail converges upon the cluster.   Actually, most models (e.g., Murali \\& Dubinsky \\cite{Mur99}) predict that the  extra-tidal material should  continue to follow  the cluster orbit and  thus take  the shape  of an  elongated  tail, or  a stream. The  stream  has been  already   revealed in dwarf  spheroidal galaxies  of  the local  group  (Mateo et al. \\cite{Mat98}), but whether the stream can also  be visible in   significantly smaller objects like  globular clusters is currently a moot point.  \\item  By constructing   the surface density    map and  performing  a Gaussian smoothing,  the low-frequency features  are enhanced over the background.  We find some marginal evidence  for a possible elongation in the extra-tidal extention based on a visual inspection of this map. This elongation  may be aligned  in a direction perdendicular to  the Galactic center, although  we already know  that the significance of this result is low; additional observations will be required to settle the  issue.    A similar  displacement  is    described  in  Fig.~3 of Johnston (\\cite{Jon98}).  \\end{enumerate}  Finally, we  want to stress the power  of the DPOSS material in  conducting this kind of  programs, either by using the standard output catalogs, as they come out  from the processing  pipeline, or the specific  re-analysis of the digitized  plate scans.  In the future we will extend this study to most of the globular clusters present on the DPOSS plates."
    },
    "0002/astro-ph0002294_arXiv.txt": {
        "abstract": "We present results from an empirical study of the Mg~II h \\& k emission lines of selected Mira variable stars, using spectra from the {\\em International Ultraviolet Explorer} (IUE).  The stars all exhibit similar Mg~II behavior during the course of their pulsation cycles.  The Mg~II flux always peaks after optical maximum near pulsation phase $\\phi=0.2-0.5$, although the Mg~II flux can vary greatly from one cycle to the next. The lines are highly blueshifted, with the magnitude of the blueshift decreasing with phase.  The widths of the Mg~II lines are also phase-dependent, decreasing from about 70 km~s$^{-1}$ to 40 km~s$^{-1}$ between $\\phi=0.2$ and $\\phi=0.6$.  We also study other UV emission lines apparent in the IUE spectra, most of them Fe~II lines. These lines are much narrower and not nearly as blueshifted as the Mg~II lines.  They exhibit the same phase-dependent flux behavior as Mg~II, but they do not show similar velocity or width variations. ",
        "introduction": "At some point in their lives, many if not most stars go through an unstable phase which leads to pulsation.  There are many classes of these pulsating stars.  Perhaps the most famous are the Cepheid variables, which are popular mostly because of their well-defined relationship between stellar luminosity and pulsation period (typically 5--50 days) that makes these stars very useful as distance indicators.  Mira variables are another important class of stellar pulsators, having long periods of 150--500 days and luminosities that vary by as much as 6--7 magnitudes from minimum to maximum.  Miras are asymptotic giant branch (AGB) stars with masses similar to that of the Sun. They have very massive, slow, cool winds, which produce a complex circumstellar environment.  Observations of molecular CO lines show that Miras are often surrounded by molecular envelopes thousands of AU in diameter.  These observations yield estimates of the wind termination velocity and total mass loss rate, which are typically of order 5 km~s$^{-1}$ and $10^{-7}$ M$_{\\odot}$ yr$^{-1}$, respectively \\citep{ky95}. The circumstellar envelopes are rich sites for dust formation and are often found to be sources of SiO, OH, and H$_{2}$O maser emission \\citep{pjd94,bl97}.  The massive winds of Miras are believed to be driven by a combination of shocks induced by stellar pulsation, and dust formation \\citep{ghb88}.  The shocks lift a substantial amount of material up to 1--2 stellar radii above the surface of the star.  Radiation pressure on dust formed in this material then pushes it away from the star. The pulsation-induced shocks not only assist in generating the massive winds of Miras, but they also determine the atmospheric structure of these stars to a large extent.  Thus, understanding the nature of the shocks and measuring their properties is essential to understanding the physics of pulsation and mass loss from pulsating stars.  The ultraviolet spectral regime is an ideal place to study radiation from the shocks.  Many UV emission lines are generated from immediately behind the shocks, which are potentially very useful diagnostics for various characteristics of the shocks.  Foremost among these lines are the strong Mg~II h \\& k lines at 2800 \\AA.  A large number of UV spectra of Miras have been taken by the {\\it International Ultraviolet Explorer} (IUE) over the years, and some of the basic characteristics of the Mg~II h \\& k lines have been noted.  It is known, for example, that the Mg~II lines are not visible throughout part of the pulsation cycle.  They typically appear at about pulsation phase $\\phi=0.1$, well after optical maximum ($\\phi=0.0$).  The Mg~II fluxes peak around $\\phi=0.3-0.45$ and then decrease until becoming undetectable at about $\\phi=0.7$ \\citep{ewb86,dgl96}.  For a set of LW-HI observations of S~Car and R~Car, \\citet{jab89} showed that the Mg~II h \\& k lines are blueshifted relative to the stellar rest frame by as much as 100 km~s$^{-1}$, and the h line is significantly stronger than the k line.  Both of these properties are very difficult to explain, as the shock speeds should be much lower than 100 km~s$^{-1}$, and for other astronomical targets the k line is almost always found to be stronger than the h line \\citep[e.g.][]{rdr95}.  Clearly the unusual behavior of the Mg~II lines of Miras should be looked at in more detail to understand the pulsation process.  In this paper, we utilize the extensive IUE data sets that exist for 5 Miras to fully characterize the behavior of the ultraviolet emission of these stars.  Many pulsation cycles are sampled for each star, allowing us to see how the UV emission lines behave from one cycle to the next. ",
        "conclusions": "We have compiled IUE observations of 5 Mira variables with substantial IUE data sets in order to study the properties of emission lines seen in the UV spectra of these stars, which are believed to be formed behind outwardly propagating shocks in the atmospheres of these pulsating stars.  Our findings are summarized as follows: \\begin{description} \\item[1.] We confirm the phase-dependent Mg~II flux behavior previously reported for Mira variables \\citep*[e.g.][]{ewb86}, which is observed for all the pulsation cycles that we study:  the Mg~II flux rises after optical maximum, peaks near $\\phi=0.2-0.5$, and then decreases.  For some Miras (e.g.\\ R~Car) the amount of Mg~II flux produced during a pulsation cycle can vary by 2--3 orders of magnitude from one cycle to the next, while for others (e.g.\\ S~Car) the flux behavior is more consistent. \\item[2.] The Mg~II k lines are almost always contaminated with circumstellar absorption lines of Fe~I and Mn~I, making analysis of the line profile very difficult. \\item[3.] The Mg~II h line is always blueshifted, with the magnitude of the blueshift decreasing with pulsation phase.  The blueshifts vary somewhat from star to star and cycle to cycle, but typical velocity changes are from $-70$ km~s$^{-1}$ to $-40$ km~s$^{-1}$ from $\\phi=0.2$ to $\\phi=0.6$.  Note, however, that these line shifts do not represent the actual gas velocities at the formation depths of these lines, because of the high opacity of Mg~II h \\& k.  These velocity variations are very similar to those of the optical Ca~II H \\& K lines. \\item[4.] The width of the Mg~II h line decreases from about 70 km~s$^{-1}$ to 40 km~s$^{-1}$ between $\\phi=0.2$ and $\\phi=0.6$. \\item[5.] In addition to the Mg~II lines, other lines of Fe~II, Fe~I, and Al~II] are also observed in IUE LW-HI spectra.  The fluxes of these lines show the same phase-dependent behavior as the Mg~II lines. \\item[6.] Unlike Mg~II, these other emission lines tend to be very narrow and do not show phase-dependent velocity and width variations.  Except for Fe~II $\\lambda$2599, the Fe~II and Al~II] lines of most of the Miras show blueshifts of $5-15$ km~s$^{-1}$, which may indicate the flow velocity of the shocked material.  In contrast, the lines of T~Cep do not show any significant line shifts, although we speculate that perhaps this is due to an uncertain center-of-mass velocity for this star. \\end{description}"
    },
    "0002/astro-ph0002157_arXiv.txt": {
        "abstract": "We report the detection of $\\sim$1\\,500 RR Lyrae of Bailey type ab located in the Sagittarius dwarf galaxy (Sgr). These variables have been detected on two ESO Schmidt fields centred on (l,b)=(3.1\\de,-7.1\\de) and (6.6\\de,-10.8\\de), covering an area of $\\sim$50 deg$^{2}$. We present a surface density map of Sgr based on the spatial distribution of these RRab, allowing us to trace its structure in a region that was still almost unexplored between b=-14\\de and b=-4\\de. We present the results of the fit of different models to the density profile of Sgr. The best fit to the core of Sgr is an exponential with a scale length of 4.1\\de along the major axis. When we look at the extension of Sgr we find a break (significant at the $\\sim$2$\\sigma$ level) in the slope of the surface density along the main axis of Sgr. The nearly flat (or at least very slowly decreasing) profile in the outer region of Sgr shows that this dwarf galaxy is probably extending even further out our fields. ",
        "introduction": "The Sagittarius dwarf galaxy is the closest known member of the Local Group orbiting around the Milky Way ($\\sim$25 kpc from the sun, $\\sim$16 kpc from the Galactic Centre), but as a consequence of its location behind the Galactic Centre, it has been discovered only recently (Ibata, Gilmore, Irwin 1994, 1995). Since this discovery it turned out that Sgr presents typical features of a dwarf spheroidal: domination of an old ($\\gtrsim$10 Gyr) metal poor stellar population (Mateo et al. \\cite{muskkk}; Fahlman et al. \\cite{fahlman}; Marconi et al. \\cite{marconi}; Bellazzini et al \\cite{bfb1}) and absence of gas (Burton \\& Lockman \\cite{bl}). Its highest surface density region is centred on the Globular Cluster M54 (l=5.6\\de, b=-14.0\\de) and it is oriented roughly perpendicular to the Galactic plane so that its Northern extension (in Galactic coordinates) is completely hidden by the MW.\\\\ The mapping of Sgr is difficult to achieve because of the combination of its low surface brightness ($\\mu_{V}\\ge 25.5$ mag.arcsec$^{-2}$), contamination by foreground Galactic stars and its large spatial extent (at least 22$^{\\circ}\\times$8\\de) (Ibata et al. \\cite{iwgis}, hereafter IWGIS). Evidence for the presence of Sgr has been established over 45\\de from b$\\sim -3^{\\circ}$ (Alard \\cite{a96}, hereafter A96; Alcock et al. 1997, hereafter Alc97) down to b$\\sim -48^{\\circ}$ (Mateo et al. \\cite{mom}, hereafter MOM), but it is difficult to assess whether these regions still correspond to the main body of Sgr or if we are merely encountering tidal debris (as suggested by Johnston et al. \\cite{johnston99}). IWGIS proposed a map of the Southern part of Sgr based on the spatial distribution of the bright main sequence stars in Sgr and covering an area of $\\sim 150$ deg$^{2}$ from $b\\sim -11^{\\circ}$ down to $b\\sim -26^{\\circ}$. However, their method based on statistical decontamination fails at low Galactic latitudes ($|$b$|\\lesssim$12\\de) where differential reddening and high density of foreground stars (only $\\sim$1 star in 1\\,000 is in Sgr in these regions) preclude any reliable decontamination, leaving the structure of the Northern extension of Sgr almost unknown. To this point, the detection of RR Lyrae constitutes an essential tool to trace the structure of Sgr in these regions as they can be clearly separated from the RR Lyrae of the MW. This method has already proven successful and $\\sim 350$ RRab were detected between b=-10\\de and b=-4\\de (A96; Alc97). However, a connection between these stars and the centre of Sgr was necessary in order to offer a clear vision of this important region strongly interacting with the MW.\\\\ In this paper we report the detection of $\\sim$1\\,500 RRab members of Sgr and located in its Northern extension. We present a surface density map of Sgr covering $\\sim$50 deg$^{2}$ between b=-14\\de and b=-4\\de, based on the spatial distribution of these variables. \\\\ The paper is organized as follows : in section 2 we present our data (observations and reduction). Section 3 is devoted to the description of the selection process of RR Lyrae stars as well as a study of its completeness. We then describe the structure of Sgr (section 4). Finally we summarize our results and conclude in section 5. ",
        "conclusions": "To summarize, we presented the detection of $\\sim$1\\,500 RRab stars located in the Sgr dwarf galaxy. A surface density map based on the spatial distribution of these variables unveiled the structure of this dwarf galaxy in a region that was still almost unexplored so far between b=-14\\de and b=-4\\de. The core of Sgr is best fitted by an exponential with a scale length of 4.1\\de along the major axis. A cross section of this density map revealed a break in the slope occurring at $\\sim$6\\de from the highest density region of Sgr and an almost flat density past the break.\\\\ Although the break coincided with the change of field we have shown that this is unlikely to be an experimental effect since it is also perceptible in the uncorrected density, whereas the \\duo field is intrinsically more sensitive to crowding than \\sag (lower resolution, lower extinction). Also, as shown in Section 4.2, the amplitude cuts used in this study cannot be considered as responsible for the break. Finally, could this break be a consequence of an overestimation of the completeness correction in \\duo relative to \\sag ?  Though not excluded, this would be in conflict with what is observed in the overlap where 3 RRab blended by a neighbouring star were detected in \\sag and missed in \\duo, a result that is quite consistent with the corrections actually applied. We argue therefore that the break is real. The significance of the break relative to an exponential with a scale length of 4.1\\de is $\\sim$2$\\sigma$. MOM also observed a break in their density profile in the Southern extension of Sgr. However neither the location (20\\de from the centre) nor the density at the break location ($\\Sigma_{V}\\sim$29.0 mag.arcsec$^{-2}$) are consistent with our values (6\\de and $\\sim$26.7 mag.arcsec$^{-2}$) implying that either the main body of Sgr is not symmetric or these ``post-break'' stars are not directly related to it.\\\\ Another striking feature revealed by the surface density profile is its flatness past the break. This feature relies on the accuracy of the completeness correction over the field, a correction that becomes quite important at low Galactic latitudes (up to 60$\\%$). Yet, the difficulty of modeling point spread functions on photographic plates (due to non-linear response of the emulsion) and potential systematic errors caused by differential sensitivity over the plate makes the crowding correction rather uncertain. Therefore, although our completeness corrections are fairly consistent within the overlap, we cannot exclude that the flatness of the density profile in the outer regions is a consequence of an overcorrection. Wide-field high resolution imaging would be necessary in these extremely crowded regions (up to $\\sim$10$^{6}$ stars per square degree at our magnitude limit) to confirm or to rule out this issue. Nevertheless, even if we consider that our completeness corrections are overestimated by a factor of 2 (a quite conservative estimate), it remains that the density profile decreases slowly in the outer regions and Sgr may well be extending even further out towards (beyond ?) the Galactic plane. \\\\ Johnston et al. (\\cite{johnston99}) recently modeled the Sgr stream as a superposition of a main body and tidal streams of stars stripped on previous peri-centric passages. This scenario has been worked out to explain both the break observed by MOM and the possible detection of stars in the outer region of Sgr with different radial velocities relative to those of the main body (Majewski et al. \\cite{maj}). Similarly, spectroscopic observations on our RR Lyrae catalogue could allow to determine the nature of the stars in the outer region: if these stars are linked to the main body of Sgr, then they should share almost the same radial velocities as the main body (apart of a gradient along the main axis due to the rapidly varying Galactic potential). On the other hand, if the break we observe corresponds to a transition between the main body and an unbound tidal stream from a previous orbit, it is likely that the two objects will have different radial velocities. This new catalogue of RR Lyrae is an interesting opportunity to study further a region of Sgr that has been poorly investigated so far.\\\\"
    },
    "0002/astro-ph0002227_arXiv.txt": {
        "abstract": "We present sub-arcsecond thermal infrared imaging of HD~98800, a young quadruple system composed of a pair of low-mass spectroscopic binaries separated by 0.8$''$ (38 AU), each with a K-dwarf primary. Images at wavelengths ranging from 5 to 24.5 $\\mu$m show unequivocally that the optically fainter binary, HD~98800B, is the sole source of a comparatively large infrared excess upon which a silicate emission feature is superposed. The excess is detected only at wavelengths of 7.9~$\\mu$m and longer, peaks at 25 $\\mu$m, and has a best-fit black-body temperature of 150~K, indicating that most of the dust lies at distances  greater than the orbital separation of the spectroscopic binary. We estimate the radial extent of the dust with a disk model that approximates radiation from the spectroscopic binary as a single source of equivalent luminosity.  Given the data, the most-likely values of disk properties in the ranges considered are $R_{in} =  {5.0}\\pm2.5$ AU, $\\Delta R = 13\\pm8$ AU, $\\lambda_0 = {2}^{+4}_{-1.5}\\mu$m, $\\gamma = 0\\pm2.5$, and $\\sigma_{total} = 16\\pm3$ AU$^2$, where $R_{in}$ is the inner radius, $\\Delta R$ is the radial extent of the disk,  $\\lambda_0$ is the effective grain size, $\\gamma$ is the radial power-law exponent of the optical depth, $\\tau$, and $\\sigma_{total}$ is the total cross-section of the grains. The range of implied disk masses is 0.001--0.1 times that of the moon. These results show that, for a wide range of possible disk properties, a circumbinary disk is far more likely than a narrow ring. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002011_arXiv.txt": {
        "abstract": "Gravitational lensing magnifies the observed flux of galaxies behind the lens. We use this effect to constrain the total mass in the cluster Abell 1689 by comparing the lensed luminosities of background galaxies with the luminosity function of an undistorted field. Under the assumption that these galaxies are a random sample of luminosity space, this method is not limited by clustering noise. We use photometric redshift information to estimate galaxy distance and intrinsic luminosity. Knowing the redshift distribution of the background population allows us to lift the mass/background degeneracy common to lensing analysis. In this paper we use 9 filters observed over 12 hours with the Calar Alto 3.5m telescope to determine the redshifts of 1000 galaxies in the field of Abell 1689.  Using a complete sample of 146 background galaxies we measure the cluster mass profile. We find that the total projected mass interior to $0.25\\mpcoh$ is $M_{2d}(<0.25\\mpcoh)=(0.48\\pm0.16)\\times10^{15} \\,h^{-1}{\\rm M}_{\\odot}$, where our error budget includes uncertainties from the photometric redshift determination, the uncertainty in the offset calibration and finite sampling. This result is in good agreement with that found by number count and shear--based methods and provides a new and independent method to determine cluster masses. ",
        "introduction": "The use of gravitational lensing as a means of cluster mass reconstruction provides a theoretically efficient approach without the equilibrium and symmetry assumptions which typically accompany virial and X-ray temperature methods. Mass determination through application of lens shear proves to give good resolution in mass maps although measurement of absolute quantities is not possible without external calibration. This so called sheet-mass degeneracy (Falco, Gorenstein \\& Shapiro 1985) is broken however by methods which exploit the property of lens magnification.  First recognised by Broadhurst, Taylor \\& Peacock (1995, BTP hereafter) as a viable tool for the reconstruction of cluster mass, lens magnification has the twofold effect of amplifying background source galaxy fluxes as well as their geometrical size and separation. This immediately permits two separate approaches for measuring lensing mass. The first involves selecting a sample of sources with a flat or near-flat number count slope. Magnification results in a reduction of their local surface number density owing to the dominance of their increased separation over the enhanced number detectable due to flux amplification.  Although contaminated by faint cluster members, Fort, Mellier \\& Dantel-Fort (1997) first reported this dilution effect using B and I band observations of the cluster CL0024$+$1654. Later, Taylor et al. (1998, T98 hereafter) demonstrated how the dilution in surface number density of a colour-selected sample of red galaxies lying behind the cluster Abell 1689 enables determination of its total mass profile and 2d distribution.  A projected mass interior to $0.24\\mpcoh$ of $M_{2d}(<0.24\\mpcoh)=(0.50\\pm0.09) \\times10^{15}\\,h^{-1}{\\rm M}_{\\odot}$ was predicted, in good agreement with the shear analysis of Tyson \\& Fischer (1995) who measured $M_{2d}(<0.24\\mpcoh)=(0.43\\pm0.02)\\times10^{15}\\,h^{-1} {\\rm M}_{\\odot}$ and Kaiser (1995) with a measurement of $M_{2d}(<0.24\\mpcoh)=(0.43\\pm0.04)\\times10^{15}\\,h^{-1} {\\rm M}_{\\odot}$.  Since then, several authors have detected source number count depletion due to cluster lensing. Athreya et al (1999) observe MS1008$-$1224 and use photometric redshifts to identify a background population of galaxies within which they measure depletion.  Mayen \\& Soucail (2000) constrain the mass profile of MS1008$-$1224 by comparing to simulations of depletion curves. Gray et al (2000) measure the first depletion in the near infra-red due to lensing by Abell 2219. Finally and most recently, R\\\"{o}gnvaldsson et al. (2000) find depletion in the source counts behind CL0024$+$1654 in the R band and for the first time, in the U band.  The second mass reconstruction approach permitted by magnification forms the primary focus of this paper.  The amplification of flux by lens magnification introduces a measurable shift in the luminosity function of background source galaxies. With a sufficiently well defined luminosity function derived from an unlensed offset field for comparison, this shift can be measured to allow an estimate of the lens mass (BTP).  This method relies upon a set of observed source magnitudes which, if assumed to form an effective random sampling of luminosity space, is not limited by noise from background source clustering unlike the number count method (see Section \\ref{sec_morph} for further discussion).  This paper presents the first application of mass reconstruction using lens flux magnification inferred from the luminosity function of background samples. Unlike the method of T98 who defined their background sample based on colour cuts, in this work photometric redshifts of all objects in the observed field have been estimated. This not only allows an unambiguous background source selection but alleviates the need to estimate source distances when scaling convergence to real lens mass.  The following section details the theory of mass reconstruction from lens magnification of background source magnitudes. Section \\ref{sec_photo_anal} describes the photometric analysis applied to observations of A1689 with the redshifts which result. Observations of the offset field which provide the absolute magnitude distribution required for comparison with the A1689 background source sample are presented in Section \\ref{sec_cadis_field}.  From this, a parameterised luminosity function is calculated in Section \\ref{sec_cadis_schechter} necessary for application of the maximum likelihood method. Following a discussion of sample incompleteness in Section \\ref{sec_completeness}, a mass measurement of A1689 is given in Section \\ref{sec_mass_determ} where the effects of sample incompleteness are quantified. Finally, a signal to noise study is carried out in Section \\ref{sec_sn_calcs} to investigate the effects of shot noise, calibration uncertainty of the offset field and photometric redshift error. ",
        "conclusions": ""
    },
    "0002/astro-ph0002083_arXiv.txt": {
        "abstract": "We present the results of the first major systematic submillimetre survey of radio galaxies spanning the redshift range $1 < z < 5$. The primary aim of this work is to elucidate the star-formation history of this sub-class of elliptical galaxies by tracing the cosmological evolution of dust mass. Using SCUBA on the JCMT we have obtained 850-\\micron{} photometry of 47 radio galaxies to a consistent rms depth of $1\\;$mJy, and have detected dust emission in 14 cases.  The radio galaxy targets have been selected from a series of low-frequency radio surveys of increasing depth (3CRR, 6CE, etc), in order to allow us to separate the effects of increasing redshift and increasing radio power on submillimetre luminosity.  Although the dynamic range of our study is inevitably small, we find clear evidence that the typical submillimetre luminosity (and hence dust mass) of a powerful radio galaxy is a strongly increasing function of redshift; the detection rate rises from $\\simeq$15 per cent at $z < 2.5$ to $\\gtrsim$75 per cent at $z > 2.5$, and the average submillimetre luminosity rises at a rate $\\propto (1+z)^3$ out to $z \\simeq 4$.  Moreover our extensive sample allows us to argue that this behaviour is not driven by underlying correlations with other radio galaxy properties such as radio power, radio spectral index, or radio source size/age.  Although radio selection may introduce other more subtle biases, the redshift distribution of our detected objects is in fact consistent with the most recent estimates of the redshift distribution of comparably bright submillimetre sources discovered in blank field surveys. The evolution of submillimetre luminosity found here for radio galaxies may thus be representative of massive ellipticals in general. ",
        "introduction": "Although large numbers of `normal' galaxies have now been discovered out to $z \\simeq 5$ (e.g. Steidel et al. 1999\\nocite{sag99}), radio galaxies continue to offer the best opportunity to study examples of {\\em massive elliptical galaxies} (or the progenitors thereof) back to comparably early cosmic epochs. The reason for this is that a powerful radio source requires a massive black hole, and it is relatively certain that nowadays all such massive black holes reside in massive ellipticals \\cite{magorrian98,mkd99}.  Thus radio selection offers a relatively efficient way of studying the properties of a sub-set of massive ellipticals as a function of cosmic epoch. Moreover this subset may well be representative of massive ellipticals in general, especially since it can be argued that a substantial fraction of all present-day ellipticals brighter than $2L^{\\star}$ must have been active at $z \\simeq 2.5$ (Dunlop et al. 2000\\nocite{dunlopetal2000}).  At low redshift, radio galaxies are dominated by well-evolved stellar populations \\cite{nolan2000}, have rather low dust/gas masses \\cite{kp91}, and lie in the same region of the fundamental plane as normal inactive ellipticals (McLure et al. 1999; Dunlop et al. 2000\\nocite{mkd99,dunlopetal2000}).  Moreover current evidence suggests that they have evolved only passively since at least $z \\simeq 1$ \\cite{ll84,mcd00}.  This, coupled with the relatively old ages derived for a few radio galaxies at $z \\simeq 1.5$ \\cite{dunlophy}, points towards a high redshift of formation, $z > 3$, for the bulk of their stellar populations.  This means that both the star-formation rate and gas mass in these galaxies should be a strongly increasing function of redshift as one approaches their primary formation epoch(s).  As has been argued by many authors, a massive starburst at high redshift is expected to produce rapid chemical enrichment and to be largely enshrouded in dust.  Consequently , submillimetre luminosity should be a good indicator of the evolutionary state of an elliptical galaxy, being expected to peak roughly half-way through the production of a galaxy from primordial material \\cite{ee96,frayerbrown97}, or even earlier (i.e. as soon as the gas receives sufficient heating) given pre-existing enrichment (Hughes, Dunlop \\& Rawlings 1997\\nocite{hdr97}). Indeed submillimetre luminosity, viewed as a tracer of gas mass, is arguably the best way to assess the evolutionary status of a massive galaxy.  There are two reasons for this.  Firstly, the sensitivity and bandwidth limitations of present-day instruments makes detecting molecular line emission from high-redshift objects extremely difficult.  Secondly, in the case of radio galaxies, optical-UV measures may be confused by the direct or indirect effects of AGN activity.  The detectability of high-redshift radio galaxies at submillimetre wavelengths was first demonstrated by Dunlop et al. \\shortcite{dhr94}, when they detected 4C41.17, at the time the most distant known galaxy at $z=3.8$.  However the observation, made with the single element bolometer detector UKT14 on the James Clerk Maxwell Telescope (JCMT), required 4 hours of integration in exceptional weather conditions.  In fact, the sensitivity of UKT14 permitted only the most extreme objects to be detected at high redshift, and it was often a struggle to detect even those.  The advent of the Submillimetre Common-User Bolometer Array (SCUBA) on the JCMT offered the first real opportunity to rectify this situation. Although photometric observations do not fully exploit the multiplex advantage offered by an array camera, the individual bolometers in the SCUBA array offered almost an order of magnitude improvement in sensitivity over UKT14.  This made it feasible to consider undertaking the first major submillimetre study of radio galaxies spanning a wide range of redshifts, and it is the results of the first major SCUBA survey of radio galaxies which we report here.  The layout of the paper is as follows: In \\sec{samp} we give a more detailed overview of pre-existing submillimetre observations of radio galaxies, and explain the motivation for observing a sample of galaxies compiled from flux-limited radio surveys of increasing depth. The resulting sample is then described and summarized, before the results of our new submillimetre observations are presented in \\sec{submmobs}.  \\sec{synchcorrect} then gives details of the radio properties of each galaxy, and explains how the total and, where possible, core radio spectrum has been extrapolated to submillimetre wavelengths to estimate (or at least constrain) the potential level of non-thermal contamination at 850$\\;$\\micron{}. The coverage of the radio-luminosity:redshift plane provided by our observed sample is presented in \\sec{pzplanesec}, and then in \\sec{seclumin} we calculate the rest-frame 850-\\micron{} luminosities or upper limits for all the observed galaxies.  In \\sec{evolstats} we present a detailed statistical exploration of the evidence for genuine cosmological evolution of \\luminsub{} in our sample. Finally in \\sec{concsec} we conclude by considering our results in the context of the recently published blank-field submillimetre surveys.  We have deliberately confined this paper to a determination and discussion of the relative behaviour of submillimetre luminosity in our sample, and have postponed our analysis and interpretation of more model-dependent properties, such as inferred gas mass and galaxy age, to a subsequent paper.  \\subsection{Conventions}  For clarity, we summarize here a number of conventions which have been adopted throughout this paper. \\begin{enumerate} \\item For each galaxy in the sample, information has been collated from several references.  Each reference has been given a code; for example, ER93\\nocite{er93} corresponds to Eales S.A., Rawlings S., 1993, ApJ, 411, 67.  These codes are included in the bibliography. \\item Throughout the paper, upper limits are calculated using the following prescription: Consider an observation with signal S and standard error $\\varepsilon$.  If the signal is positive, the $n$-$\\sigma$ upper limit is S$+(n\\times\\varepsilon)$.  If the signal is negative, the $n$-$\\sigma$ upper limit is $n\\times\\varepsilon$.  It cannot be guaranteed that the upper limits taken from other papers were calculated in this manner. \\item For a power-law spectrum, the spectral index, $\\alpha$, is defined as ${\\rm S}_{\\nu}\\propto \\nu^{-\\alpha}$, where S$_{\\nu}$ is the flux density at frequency $\\nu$.  Thus, a radio synchrotron spectrum has a positive spectral index, and the Rayleigh-Jeans tail of a thermal spectrum has a negative spectral index. \\item The cosmological constant, $\\Lambda$, is assumed to be zero throughout the paper. \\end{enumerate} ",
        "conclusions": "\\label{concsec}  In summary, our attempts to investigate and quantify the various possible biases which might conceivably afflict this study have simply served to reaffirm and strengthen our basic result, namely that the submillimetre luminosity of radio galaxies is primarily a function of redshift as illustrated in \\Fig{ombinning}.  It therefore seems hard to avoid the straightforward conclusion that the observed increase in submillimetre detection rate and characteristic luminosity with redshift is due to the increasing youthfulness of the stellar populations of the radio galaxies in our sample.  In a separate paper we will explore how the inferred evolution of gas mass and star-formation rate in these galaxies compares with the predictions of models of elliptical galaxy formation and evolution. However, it is interesting to briefly consider whether the apparently rather extreme evolution is peculiar to radio galaxies, or may in fact be typical of the cosmological evolution of dust and gas in massive ellipticals in general.  In \\Fig{cumulrgzdistrib} we show the cumulative redshift distribution of our radio galaxy sample, along with the cumulative redshift distribution of the subset detected at submillimetre wavelengths. This figure serves to re-emphasize that the high median redshift of our detected galaxies ($z = 3$) does not simply reflect the median redshift of the sample selected for observation ($z = 2$). However, what is particularly interesting is that the redshift distribution of our submillimetre detected radio galaxies is statistically indistinguishable from current best estimates of the redshift distribution of sources detected in SCUBA surveys, as illustrated in \\Fig{smailz} \\cite{smailzdistrib}.  This comparison provides at least circumstantial evidence that the submillimetre evolution of radio galaxies found here may indeed be symptomatic of the evolution of massive elliptical galaxies in general.  At first sight, it may appear contradictory that the average submillimetre luminosity of the radio galaxy sample continues to rise beyond $z\\sim3$ while the median redshift of the sample is $z\\sim3$. This results from the median redshift of our most luminous detections being higher at $z\\sim3.6$ for sources brighter than 5$\\;$mJy.  This raises the interesting possibility that the most massive dust-enshrouded starbursts are confined to $z>3$.  Therefore the median redshift of submillimetre sources detected in blank-field surveys may prove to be a function of flux density.  It will be some time before the redshift information in bright submillimetre surveys approaches that currently available for radio-selected samples. However, it will undoubtedly be very interesting to see how this comparison evolves as submillimetre-selected samples are studied and refined in the years to come."
    },
    "0002/astro-ph0002426_arXiv.txt": {
        "abstract": "We investigate the role of the eccentric disc resonance in systems with mass ratios $q \\age 1/4$, and demonstrate the effects that changes in the mass flux from the secondary star have upon the disc radius and structure. The addition of material with low specific angular momentum to its outer edge restricts a disc radially. Should the mass flux from the secondary be reduced, it is possible for the disc in a system with mass ratio as large as $1/3$ to expand to the $3:1$ eccentric inner Lindblad  resonance and for superhumps to be excited. ",
        "introduction": "The superhump phenomenon was first observed in superoutbursts of dwarf novae of short orbital period ($P_{\\rm orb} \\ale 2.8\\,{\\rm h}$: see Warner 1995b). It can be explained as an effect of an eccentric disc that precesses with respect to the tidal field of the secondary star (Whitehurst 1988). In low mass ratio binaries ($q \\ale 1/4$), disc eccentricity can be excited via tidal resonance (Whitehurst 1988, Hirose \\& Osaki 1990, Lubow 1991). In systems with larger $q$ however, the disc is truncated at too small a radius for the resonance to be effective.  Observationally there is evidence that discs can be tidally resonant for $q$ significantly greater than $1/4$.  Superhumps appear in high mass transfer discs (those of nova-like variables and of nova remnants) in some systems with orbital periods up to 3.5 h (Skillman et al. 1998). At this period the mass of the secondary star is $0.31\\,\\msol$ (equation 2.100 of Warner 1995b). Unfortunately there are no reliable determinations of $q$ in the superhumping systems with $\\Pb > 3$h. The mean mass of the primaries of the nova-likes is $\\sim 0.74\\,\\msol$ (Warner 1995b) and the range is likely to be at least $0.5 - 1.0\\,\\msol$. If the primary mass is at the upper end of this range, it is therefore possible that a system with $\\Pb \\sim 3.5$h would have a mass ratio approaching that needed for the eccentric resonance to lie within the accretion disc.  There can be little doubt that a mechanism is required to extend the range of mass ratios for which tidal resonance occurs, beyond the $q \\ale 1/4$ established in equilibrium disc calculations. One possibility is suggested by the VY~Scl stars, systems  with orbital periods in the range $3.0 < \\Pb < 4.0$ h (Warner 1995b) in which mass transfer is liable to reduce significantly  at irregular intervals. Such a drop in $\\dot M_{\\rm s}$ would allow a tidally stable, equilibrium disc that just marginally fails to reach the eccentric resonance to expand radially and become eccentrically unstable.  In this paper we demonstrate that the discs in high $\\dot M$ systems can be excited into an eccentric state when the mass transfer from the secondary drops below its equilibrium value. In the next section we discuss the observational peculiarities of the VY~Scl stars. In section three we outline our model for superhumps in these systems. We present numerical results that support our case in section four, and make our conclusions in section five. ",
        "conclusions": "Under conditions of steady mass transfer, a mass ratio $\\ale 1/4$ is required for a close binary accretion disc to encounter the $3:1$ eccentric inner Lindblad resonance. However, it is possible for eccentricity to be excited in the disc of a high mass transfer system with $q \\ale 1/3$ if $\\dot M_{\\rm s}$ is reduced, as is thought to occur in the VY~Sculptoris systems. Our simulations suggest that the disc will remain eccentric as long as  $\\dot M_{\\rm s}$ remains at the lower value.  Thus, a precessing, eccentric disc remains the best explanation of superhumps, even for systems with $3.0 < \\Pb < 4.0$ h."
    },
    "0002/astro-ph0002176_arXiv.txt": {
        "abstract": "We adapt and modify the eigenfunction method of computing the power-law spectrum of particles accelerated at a relativistic shock front via the first-order Fermi process \\cite{aguthmann:kirkschneider87} to apply to shocks of arbitrarily high Lorentz factor. The power-law index of accelerated particles undergoing isotropic small-angle scattering at an ultrarelativistic, unmagnetized shock is found to be $s=4.23\\pm0.2$ (where $s=d\\ln f/ d\\ln p$, with $f$ the Lorentz-invariant phase-space density and $p$ the momentum), in agreement with the results of Monte-Carlo simulations. We present results for shocks in plasmas with different equations of state and for Lorentz factors ranging from 5 to infinity. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002340_arXiv.txt": {
        "abstract": "The Toronto Red-Sequence Cluster Survey (TRCS) is a new galaxy cluster survey designed to provide a large sample of optically selected $0.1 < z < 1.4$ clusters. The planned survey data is 100 square degrees of two color ($R$ and $z'$) imaging, with a 5$\\sigma$ depth $\\sim$2 mag past $M^*$ at $z=1$. The primary scientific drivers of the survey are a derivation of $\\Omega_{m}$ and $\\sigma_8$ (from $N(M,z)$ for clusters) and a study of cluster galaxy evolution with a complete sample. This paper gives a brief outline of the TRCS survey parameters and sketches the methods by which we intend to pursue the main scientific goals, including an explicit calculation of the expected survey completeness limits.  Some preliminary results from the first set of data ($\\sim$ 6 deg$^2$) are also given. These preliminary results provide new examples of rich $z\\sim1$ clusters, strong cluster lensing, and a possible filament at $z\\sim1$. ",
        "introduction": "\\paragraph{} The Toronto Red-Sequence Cluster Survey (TRCS) is a major new observational effort designed to identify and characterize a large sample of galaxy clusters to redshifts as high as $z\\sim1.4$. When completed, the TRCS will be the largest imaging survey ever completed on 4m telescopes, and will provide a large and homogeneous sample of galaxy clusters for detailed follow-up study. The basic survey is envisioned as 100 deg$^2$ of 2 filter ($R$ and $z'$) imaging, to a depth which is $\\sim$2 mag past $M^*$ at $z=1$ in both filters. The design of the survey is based on a new method for identifying galaxy clusters (Gladders \\& Yee  2000a) developed specifically for the TRCS. In brief, this method searches for clustering in the 5-D space of: x-y positions, $R-z'$ color, $z'$ mag, and morphology in the form of a concentration index. The x-y positions provide the surface density enhancement. A color slice in the color-mag plane provides separation in $z$ space via the {\\it red sequence} of early-type galaxies in clusters (Figure 1) and increases the S/N of density enhancements. Morphology allows us to key onto early-type galaxies, the primary population in cluster centers.  \\begin{figure} \\plotone{gladders1.eps} \\caption{ Modeled cluster CMDs to z=1.4 (left panel, solid lines), from Kodama \\& Arimoto (1997).  Diamonds indicate simulated field galaxies for a 1 arcmin$^2$ FOV. The *s show M$^{*}$ for each redshift. The TRCS photometric completeness limits are shown (dashed line). We also show real CMDs (right) for CL1322+3114 at $z=0.75$ and k-corrected to $z=1$. These data are from HST images degraded to TRCS seeing and depth.  Cluster image objects (*) and field image objects ($\\diamond$) are shown. Note the visibility of the cluster red sequence.} \\end{figure}  \\subsection{Scientific Goals} \\paragraph{} The TRCS is being driven by two major scientific goals. The first is based on the theoretical prediction that the evolution of the mass-spectrum of galaxy clusters with redshift, $N(M,z)$, should be a strong function of two cosmological parameters, $\\Omega_{m}$ and $\\sigma_8$ (Figure 2). The goal is to use the clusters identified in the survey to measure $N(M,z)$ directly from the survey data. Redshift can be estimated from the color of red sequence (e.g., L\\'{o}pez-Cruz \\& Yee 2000), and the mass of each cluster can be estimated from its richness, as measured by the parameter $B_{gc}$ (e.g., Yee \\& L\\'{o}pez-Cruz 1999). The second major scientific goal is a study of the cluster galaxy populations, which can be done using the TRCS for the first time with a complete sample. The definition of a complete, or volume limited, sample is derived from extensive simulations of the survey selection functions.  \\begin{figure}[htp] \\plotfiddle{gladders2.eps}{3cm}{0.0}{50}{40}{-165}{-40} \\caption{The expected cumulative counts of clusters per~deg$^2$ for two cosmologies, for Abell Richness Class (ARC) 1 and 2.  The 100\\% completeness redshift for ARC 1 is $\\sim1.1$, and $\\sim1.3$ for the richer, rarer ARC 2 clusters.} \\end{figure}   \\subsection{Survey Completeness and Selection Functions} \\paragraph{} Any detailed understanding of the cosmological or galaxy evolution results deriving from the TRCS requires a good understanding of the survey selection functions. Specifically, we wish to know how well the cluster-finding algorithm finds clusters of various sorts (as described by various parameters). To this end, we have constructed a number of cluster and field simulations (Gladders \\& Yee 2000b) to directly test the algorithm. A large suite of possible clusters have been tested; the parameters describing the clusters are given in Table 1. The results of this process demonstrate that the TRCS should be complete for all reasonable clusters of Abell Richness Class $\\geq$ 1 clusters ($\\sigma_{v}\\geq750$ km s$^{-1}$) to at least $z=1.1$. \\begin{table} \\begin{tabular}{lll} \\hline Parameter & Model Values & Notes \\\\\\hline LF $R$-band $M^*$ & -22.5, -22.25, -22.0 & $\\alpha=-1.0$\\\\ Abell Richness counts & 35,44,56,72,93,120& Richness Classes 0-2 \\\\ NFW core scale radius & 0.1,0.2,0.3,0.4,0.5 & in $h^{-1}$ Mpc\\\\ ellipticity &0.0,0.2,0.4,0.6,0.8 & measured at 1 $h^{-1}$ Mpc \\\\ blue fraction & 0.1,0.5,0.65,0.8,0.9 &  \\\\ red sequence age & 9,10,11,11.5 & lower limit of SF in Gyr  \\\\ scatter in formation ages& 0.5,1.0,max & tophat width in Gyr  \\\\ cluster redshift&0--1.4&\\\\\\hline \\end{tabular} \\caption{Cluster model parameters used to test the cluster finding algorithm as applied to the TRCS.} \\end{table} ",
        "conclusions": ""
    },
    "0002/astro-ph0002030_arXiv.txt": {
        "abstract": "We provide a pedagogical overview of defect models of structure formation. We first introduce the concept of topological defect, and describe how to classify them. We then show how defects might be produced in phase transitions in the Early Universe and approach non-pathological scaling solutions. A very heuristic account of structure formation with defects is then provided, following which we introduce the tool box required for high precision calculations of CMB and LSS power spectra in these theories. The decomposition into scalar vector and tensor modes is reviewed, and then we introduce the concept of unequal-time correlator. We use isotropy and causality to constrain the form of these correlators. We finally show how these correlators may be decomposed into eigenmodes, thereby reducing a defect problem to a series of ``inflation'' problems. We conclude with a short description of results in these theories and how they fare against observations. We finally describe yet another application of topological defects in cosmology: baryogenesis. ",
        "introduction": "Phase transitions are ubiquitous in nature. Typically, as the temperature of a system drops below the critical temperature, the system makes a transition from a state with greater symmetry to one with less symmetry. In general, the state with less symmetry is not unique, but can lie anywhere in a so-called {\\it vacuum manifold}. Depending on the topology of this vacuum manifold, defects will form during the phase transition. If the topology of the vacuum manifold admits defects, then these defects will inevitably arise during the phase transition unless the dynamics is completely adiabatic.  In the early Universe the temperature was decreasing very rapidly. On length scales larger than the Hubble radius, causality prevents the system from maintaining adiabaticity through interactions, and therefore on these scales defects will arise in any cosmological phase transition in which they are topologically allowed. The Universe has undergone several phase transitions. We are quite confident about those which occurred at lower temperatures: the confinement transition at a temperature $T \\sim 10^2 {\\rm GeV}$ and the electroweak symmetry breaking transition at $T \\sim 10^3 {\\rm GeV}$. Unified field theories of fundamental interactions predict the existence of other transitions at higher temperatures, e.g. a phase transition at $T \\sim 10^{16} {\\rm GeV}$ in Grand Unified Models, the supersymmetry breaking phase transition in supersymmetric models, and various compactification transitions in string (and M-) theory.  Since topological defects carry energy density, they will curve space-time and can thus act as the seeds for gravitational accretion (see e.g. Refs. \\cite{TK80,Vil85,ShellVil,HK95,RB94} for comprehensive reviews). Since inside of topological defects the symmetry characteristic of the high temperature phase is unbroken, topological defects can interact in various interesting ways with the surrounding matter and can have an effect on cosmological issues such as baryogenesis (see e.g. \\cite{BDH,BD,BDPT}), magnetic field generation (see e.g. \\cite{DD,BZ99}), and ultra-high-energy cosmic ray production \\cite{MB,PB,Sigl,BV}.  In these lectures, we first review the classification of topological defects and explain why in models with the appropriate topology, defects will inevitably form during the symmetry breaking phase transition. In Section 3 we discuss some initial applications of topological defects to cosmology. We review the domain wall and monopole problems and explain why the cosmic string model yields a promising mechanism for structure formation.  In the following sections we provide a more technical description of how high accuracy calculations of structure formation in defect theories are performed.  We first describe the details of the scalar, vector and tensor decomposition (Section \\ref{svt}). This is an invaluable tool in linear perturbation theory. Then in Section \\ref{corrs} we introduce the concept of UETC and show their general form, assuming isotropy and scaling, but not energy conservation. In Section~\\ref{caus} we show how causality limits further the form of the correlators, in the large wavelength limit. These results will be important when checking upon the numerics.  Then in Section~\\ref{result} we present the UETCs measured for cosmic strings, and highlight some of their features. Their remarkable novelty is the dominance of the energy density over any other components of the stress energy tensor. This property sets strings apart, resulting in a dominance of scalar modes over vector and tensor modes. We present some conclusions on string scenarios of structure formation, and also hybrid scenarios combining strings and inflation.  In the final section we illustrate one application of topological defects to cosmology which involves microscopic physics rather than gravitational accretion: we discuss the basics of defect-mediated baryogenesis.    The last decade has witnessed unprecedented progress in mapping the cosmic microwave background (CMB) temperature anisotropy and the large scale structure (LSS) of the Universe. The prospect of fast improving data has forced theorists to new standards of precision in computing observable quantities. The new standards have been met in theories based on cosmic inflation\\cite{hsselj,hw}. Topological defect scenarios \\cite{ShellVil,HK95} have been more challenging. However, recently there have been a number of computational breakthroughs in defect theories, partly related to improvements in computer technology. Most strikingly, the method described in \\cite{pst} showed how one could glean from defect simulations all the information required to compute accurately CMB and LSS power spectra. In this method the simulations are used uniquely for  evaluating the two point functions (known as unequal time correlators, or UETCs) of the defects' stress-energy tensor. UETCs are all that is required for computing CMB and LSS power spectra. Furthermore, they are constrained by requirements of self-similarity (or scaling) and causality, which enable us to radically extend the dynamical range of simulations, a fact central to the success of the method.   This method was applied to theories based on global symmetries. In recent work \\cite{chm,chm1} we have shown how the same method could be applied to local cosmic strings (see also \\cite{steb,abr}). In the next few Sections we shall describe in detail the simulation and  measurement of UETCs which led to the work in \\cite{chm,chm1}; as well as present the analytical tools used in this enterprise. The formalism we had to use is unfortunately more complicated than \\cite{pst}. Local strings have an extra complication over global defects, which stems from the fact that we are unable to simulate the underlying field theory. Instead, we approximate the true dynamics with line-like relativistic strings. This is thought to be reasonable for the large scale properties of the stress-energy tensor, but we do not have a good understanding of how the string network loses energy in order to maintain scaling.  This leads to two problems. Firstly one is forced to make assumptions about which cosmological fluids pick up this deficit. It is often assumed that all the strings' energy and momentum is radiated into gravitational waves, approximated by a relativistic fluid. This is by no means certain, and it may well be that the energy and momentum is transferred to particles \\cite{VinAntHin98}, and hence to the baryon, photon and CDM components.  Secondly, it is not enough to find correlators for a reduced number of stress energy tensor components (two scalar, two vector, two tensor), and then find the others by means of energy conservation. If energy conservation can  be used then one needs to compute 3 scalar, 1 vector, and 1 tensor UETCs. If one is not allowed to make use of energy conservation, one has to compute 10 scalar, 3 vector, and 1 tensor UETC.  In the following sections we first give a qualitative description of the technical novelties introduced in defect scenarios. We describe defects as active incoherent perturbations. We then describe a set of tools with which we can perform high accuracy calculations of structure power spectra in these scenarios. ",
        "conclusions": ""
    },
    "0002/astro-ph0002206_arXiv.txt": {
        "abstract": "There are several Seyfert galaxies for which there is a discrepancy between the small column of neutral hydrogen deduced from X-ray observations and the much greater column derived from the reddening of the optical/UV emission lines and continuum. The standard paradigm has the dust within the highly ionized gas which produces O~VII and O~VIII absorption edges (i.e., a ``dusty warm absorber''). We present an alternative model in which the dust exists in a component of gas in which hydrogen has been stripped, but which is at too low an ionization state to possess significant columns of O~VII and O~VIII (i.e, a ``lukewarm absorber''). The lukewarm absorber is at sufficient radial distance to encompass much of the narrow emission-line region, and thus accounts for the narrow-line reddening, unlike the dusty warm absorber. We test the model by using a combination of photoionization models and absorption edge fits to analyze the combined {\\it ROSAT}/{\\it ASCA} dataset for the Seyfert 1.5 galaxy, NGC 3227. We show that the data are well fit by a combination of the lukewarm absorber and a more highly ionized component similar to that suggested in earlier studies. We predict that the lukewarm absorber will produce strong UV absorption lines of N~V, C~IV, Si~IV and Mg~II. Finally, these results illustrate that singly ionized helium is an important, and often overlooked, source of opacity in the soft X-ray band (100 - 500 eV). ",
        "introduction": "The presence of absorption edges of O~VII and O~VIII in the X-ray (Reynolds 1997; George et al. 1998a) indicates that there is a significant amount of intrinsic ionized material along our line-of-sight to the nucleus in a large fraction ($\\sim$ 0.5) of Seyfert 1 galaxies. In addition to highly ionized gas (referred to as an X-ray or ``warm'' absorber), X-ray spectra often show evidence for a less-ionized absorber. This component has been modeled using neutral gas (cf. George et al. 1998b), and its relationship to the ``warm'' absorber is unclear. Interestingly, there are several instances in which this additional neutral column is too small by as much as an order of magnitude to explain the reddening of the continuum and emission lines, assuming typical Galactic dust/gas ratios (cf. Shull \\& van Steenburg 1985). This inconsistency was first noted in regard to the absence of high ionization emission lines in the {\\it IUE} spectra of MCG -6-30-15 (Reynolds \\& Fabian 1995). The first quantitative comparison of the neutral columns inferred from the X-ray data to that derived from the reddening was for the QSO IRAS 13349+2438 by Brandt, Fabian, \\& Pounds (1996), who suggested that the dust exists within the highly ionized X-ray absorber (ergo, a dusty warm absorber). It has been suggested that dusty warm absorbers are present in several other Seyferts (NGC 3227: Komossa \\& Fink 1997a; NGC 3786: Komossa \\& Fink 1997b; IRAS 17020+4544: Leighly et al. 1997, Komossa \\& Bade 1998; MCG -6-30-15: Reynolds et al. 1997). Since it is unlikely that dust could form within the highly ionized gas responsible for the O~VII and O~VIII absorption, it has been suggested that the dust is evaporated off the putative molecular torus (at $\\sim$ 1 pc) and, subsequently, swept up in an radially outflowing wind (cf. Reynolds 1997).   In this paper, we present an alternative explanation. It is possible that there is a component of dusty gas (which we will refer to as the ``lukewarm'' absorber), with an ionization state such that hydrogen is nearly completely ionized but the O~VII and O~VIII columns are negligible, which has a sufficient total column to account for the reddening. Such a possibility has been mentioned by Brandt et al. (1996), while Reynolds et al. (1997) have suggested that the dusty warm absorber in MCG -6-30-15 may have multiple zones. Here we suggest that the lukewarm absorber lies far into the narrow-line region (NLR). Such a component has been detected in the Seyfert galaxy NGC 4151 (Kraemer et al. 1999), and it lies at sufficient radial distance to cover much of the NLR. We will demonstrate that the combination of a dusty lukewarm absorber and a more highly ionized (O~VII and O~VIII) absorber is consistent with the observed X-ray data and with the reddening of the narrow emission lines in the Seyfert 1 galaxy NGC 3227. ",
        "conclusions": "We have shown that the X-ray spectrum of NGC~3227 is consistent with attenuation by the sum of a highly-ionized absorber and a lukewarm absorber. We suggest that these are physically different components of the circumnuclear material surrounding NGC~3227. The characteristics of the  highly-ionized absorber are similar to those previously suggested for NGC 3227 (George et al. 1998b); this is in the range of and generally similar to those in other Seyfert 1s (Reynolds 1997; George et al. 1998a). Such absorbers have been observed to vary on timescales $\\lesssim$3~yr (and much faster in some cases), and are probably due to gas well within the NLR.  The main result of this paper is that the second component, our ``lukewarm absorber'', has the appropriate physical conditions to simultaneously explain the absorption seen below 0.5 keV in the X-ray band (previously modeled as completely neutral gas) {\\it and} the reddening seen in the optical/UV. Agreement with the soft X-ray data is the result of the lukewarm gas containing significant opacity due to He~II. Agreement with the reddening of the narrow emission lines places the component outside the NLR.   Although the lukewarm absorber has the appropriate physical conditions and radial distance to redden the NLR, it must also have a sufficiently high covering fraction to be detected. For example, Reynolds (1997) found that 4/20 of radio-quiet active galaxies show both intrinsic X-ray absorption and reddening. Thus, the global covering factor of the dusty ionized absorber must be 20\\%, within the solid angle that we see these objects (cf., Antonucci 1993). Kraemer et al. (1999) have shown that the covering factor for optically thin gas in NGC 4151, similar to our lukewarm model, can be quite large ($\\sim$ 30\\%). In addition, Crenshaw et al. (1999) find that $\\sim$ 60\\% of Seyfert 1 galaxies have UV absorbers with a global covering factor $\\geq$ 50\\%, and an ionization parameter similar to the lukewarm absorber, but with lower columns on average (cf. Crenshaw \\& Kraemer 1999). Therefore, it is entirely plausible that there would be optically thin NLR gas along our line-of-sight to the nucleus in a fraction of Seyfert 1s.  The lukewarm model predicts a column of Mg~II of 3.3 x 10$^{14}$ cm$^{-2}$, which would produce strong Mg~II $\\lambda$2800 absorption. It is interesting that NGC 3227 is one of the few Seyfert 1s to show evidence of Mg~II $\\lambda$2800 in absorption (Ulrich 1988). While Kriss (1998) has shown that Mg~II absorption can arise in clouds characterized by small column density and low ionization parameter (N$_{H}$ $\\sim$ 10$^{19.5}$ cm$^{-2}$, U $\\sim$ 10$^{-2.5}$), our results predict that it may also arise in a large column of highly ionized NLR gas, even if a substantial fraction of Mg is depleted onto dust grains.  The lukewarm model predicts average grain temperatures of 30K -- 60K, for grains with radii from 0.25 $\\mu$m -- 0.005 $\\mu$m, respectively. The reradiated IR continuum, which is produced primarily by the silicate grains (cf. Mezger, Mathis, \\& Panagia 1982), peaks near 60 $\\mu$m. Assuming a covering factor of unity, this component only accounts for $\\sim$ 1\\% of the observed IR flux from NGC 3227, (which is $\\approx$ 7.98 Jy at 60 $\\mu$m; {\\it The IRAS Point Source Catalog} [1985]). It is likely that most of the thermal IR emission in NGC 3227 arises in the dense (n$_{H}$ $\\geq$ 10$^{3}$ cm$^{-3}$) NLR gas in which the narrow emission lines are formed, as is the case for the Seyfert 2 galaxy, Mrk 3 (Kraemer \\& Harrington 1986)."
    },
    "0002/astro-ph0002212_arXiv.txt": {
        "abstract": "We analyze the wind generated by the great 20 year long super-Eddington outburst of $\\eta$-Carinae.  We show that using classical stellar atmospheres and winds theory, it is impossible to construct a consistent wind model in which a sufficiently {\\em small} amount of mass, like the one observed, is shed.  One expects the super-Eddington luminosity to drive a thick wind with a mass loss rate substantially higher than the observed one.  The easiest way to resolve the inconsistency is if we alleviate the implicit notion that atmospheres are homogeneous.  An inhomogeneous atmosphere, or ``porous\", allows more radiation to escape while exerting a smaller average force.  Consequently, such an atmosphere yields a considerably lower mass loss rate for the same total luminosity. Moreover, all the applications of the Eddington Luminosity as a strict luminosity limit should be revised, or at least reanalyzed carefully.  \\vskip 0.5cm \\centerline{\\em To appear in the Astrophysical Journal Letters} \\vskip 0.5cm ",
        "introduction": "$\\eta$-Carinae is probably one of the most remarkable stellar object to have ever been documented. About 150 years ago, the star began a 20 year long giant eruption during which it radiated a supernova-like energy of roughly $3 \\times 10^{49}~ergs$~ (\\citenp{DH97}).  Throughout the eruption it also shed some $1-2~\\ms$ of material carrying approximately $6\\times 10^{48}~ergs$ as kinetic energy (\\citenp{DH97}), while expanding at a velocity of $650~km/sec$ (\\citenp{HA92}, \\citenp{C96}). $\\eta$-Carinae can therefore serve as a good laboratory for the study of atmospheres at extreme luminosity conditions.  At first glance, it appears that the star shed a large amount of material. Indeed, the inferred mass loss rate during the great eruption of $\\sim 0.1~\\ms/yr$ is significantly larger than the mass loss rate inferred for the star today ($\\lesssim 10^{-3}~\\ms/yr$, \\citenp{DH97} and references therein). However, considering that the luminosity during the great eruption is estimated to be significantly above the Eddington limit, we shall show that the star should have had a much higher mass loss rate. In fact, it should have lost during the 20 year eruption more mass than its total mass, giving rise to an obvious discrepancy.  A review of our current knowledge of $\\eta$ Car can be found in \\cite{DH97}. In section \\ref{sec:wind} we summarize how a wind solution for the star $\\eta$ Car should be constructed. Since the luminosity is very high, the effects of convection must be taken into account. In section \\ref{sec:discrepancy} we integrate the wind equations to show that no consistent solution for $\\eta$ Car exists within the possible range of observed parameters. Section \\ref{sec:bad} is devoted to possible classical solutions to the discrepancy, showing that no such possibility exists. In section \\ref{sec:good}, we show that a porous atmosphere is a simple and viable solution to the wind discrepancy. ",
        "conclusions": "\\label{sec:summary}  To summarize, the super-Eddington luminosity emitted by $\\eta$-Car should have generated a much thicker wind with a sonic point placed significantly deeper than what can be directly inferred from the observations.  A solution which lives in harmony with observations and theoretical modeling is a porous atmosphere, which allows more radiation to escape while exerting a smaller average force. It also means that the Eddington limit is not as destructive as one would a priori think it must be, even in a globally spherically symmetric case. Namely, all astrophysical analyses that employ the Eddington limit as a strict limit should be reconsidered carefully, even if they involve only unmagnetized Thomson scattering material. If $\\eta$-Carinae could have been super-Eddington for such a long duration without ``evaporating'', other systems can display a similar behavior."
    },
    "0002/astro-ph0002024_arXiv.txt": {
        "abstract": "We present the first full FIR spectrum of Centaurus A (NGC 5128) from 43 - 196.7 $\\mu$m. The data was obtained with the ISO Long Wavelength Spectrometer (LWS). We conclude that the FIR emission in a 70~\\arcsec~beam centred on the nucleus is dominated by star formation rather than AGN activity. The flux in the far-infrared  lines is $\\sim$ 1 \\% of the total FIR: the \\cii line flux is $\\sim$ 0.4 \\% FIR and the \\oi line is $\\sim$ 0.2 \\%, with the remainder arising from \\oiiinb, \\nii and \\niii lines. These are typical values for starburst galaxies.  The ratio of the \\niii / \\nii line intensities from the HII regions in the dust lane corresponds to an effective temperature, T$_{\\mathrm{eff}}$ $\\sim$ $35\\,500$ K, implying that the tip of the main sequence is headed by O8.5 stars and that the starburst is $\\sim$ 6 $\\times 10^6$ years old. This suggests that the galaxy underwent either a recent merger or a merger which triggered a series of bursts. The N/O abundance ratio is consistent with the range of $\\sim$ 0.2 - 0.3 found for Galactic HII regions.  We estimate that $<$ 5 \\% of the observed \\cii arises in the cold neutral medium (CNM) and that $\\sim$ 10 \\% arises in the warm ionized medium (WIM). The main contributors to the \\cii emission are the PDRs, which are located throughout the dust lane and in regions beyond where the bulk of the molecular material lies. On scales of $\\sim$ 1 kpc the average physical properties of the PDRs are modelled with a gas density, n $\\sim$ $10^3$ cm$^{-3}$, an incident far-UV field, G $\\sim$ $10^2$ times the local Galactic field, and a gas temperature of $\\sim$ 250 K. ",
        "introduction": "Centaurus A (NGC 5128) is the nearest (d = 3.5 Mpc; 1 \\arcsec $\\sim$17~pc, Hui et al. 1993) example of a giant elliptical galaxy associated with a powerful radio source. The large-scale radio morphology consists of twin radio lobes separated by $\\sim$ 5 degrees on the sky. The compact ($\\sim$ milliarcsecond) radio nucleus is variable and has a strong jet extending $\\sim$ 4 arcminutes towards the northeast lobe. The spectacular optical appearance is that of a giant elliptical galaxy that appears enveloped in a nearly edge on, warped dust lane. There is also a series of faint optical shells. The stellar population in the dominant elliptical structure is old, whilst that of the twisted dust lane is young, sporadically punctuated by HII regions, dust and gas (Graham 1979). The overall structure of Cen A resembles that of a recent ($< 4 \\times 10^8$ years, Tubbs 1980) merger, between a spiral and a large elliptical galaxy. The dust lane is the source of most (90 \\%) of the far-infrared luminosity (L$_{\\mathrm{FIR}} \\sim 3 \\times 10^{9}$ L$_{\\odot}$) and is thought to be re-radiated starlight from young stars in the dusty disk (Joy et al. 1988).  In Sect. 2 we describe the observations and data analysis. Sect. 3 looks at the general FIR properties and proceeds to model the HII regions and the PDRs in the dust lane. Sect. 4 summarises the results and presents our conclusions. ",
        "conclusions": "We present the first full FIR spectrum from 43 - 196.7 $\\mu$m of Cen A. We detect seven fine structure lines (see Table 2), the strongest being those generated in PDRs. At the central position, the total flux in the far-infrared  lines is $\\sim$ 1 \\% of the total FIR luminosity (L$_{43-197 \\mu m} = 3.2 \\times 10^9$ L$_{\\odot}$ for a distance of 3.5 Mpc). The \\cii line flux is $\\sim$0.4 \\% FIR and the \\oi line flux is $\\sim$ 0.2 \\% FIR. These are typical values for starburst galaxies (Lord et al. 1996). The \\oiii 52 $\\mu$m / \\oiii 88 $\\mu$m line intensity ratio is $\\sim$ 0.9, which corresponds to an electron density, n$_{\\mathrm{e}} \\sim$ 100 cm$^{-3}$ (Rubin et al. 1994). The {\\it thermal pressure} of the ionized medium in the Cen A dust lane is closer to that of starburst galaxies (n$_e \\sim$ 250 cm$^{-3}$ in M82 (Colbert et al. 1999) and M83 (Stacey et al. 1999)) than that of the Milky Way (n$_e \\sim$ 3 cm$^{-3}$ (Pettuchowski \\& Bennett 1993)).   The \\niii / \\nii line intensity ratio is $\\sim$ 1.6, giving an abundance ratio N++/N+ $\\sim$ 0.3, which corresponds to an effective temperature, T$_{\\mathrm{eff}} \\sim$ 35\\,500 K (Rubin et al. 1994). Assuming a coeval starburst, then the tip of the main sequence is headed by O8.5 stars, and the starburst is $\\sim$ 6 $\\times 10^6$ years old. If the burst in Cen A was triggered by the spiral-elliptical galaxy merger then its occurance was very recent. Alternatively, the merger triggered a series of bursts of star formation and we are witnessing the most recent activity.  We estimate that the N/O abundance ratio is $\\sim$ 0.2 in the HII regions in Cen A. This value is consistent with the range of $\\sim$ 0.2 - 0.3 found for Galactic HII regions (Rubin et al. 1988). N/O is a measure of the chemical evolution and we expect it to increase with time (c.f. the solar value of $\\sim$ 0.12).  We estimate that $\\sim$ 10 \\% of the observed \\cii arises in the WIM. The CNM contributes very little ($< 5$ \\%) \\cii emission at our beam positions. The bulk of the emission is from the PDRs.  We derive the average physical conditions for the PDRs in Cen A for the first time. There is active star formation throughout the dust lane and in regions beyond the bulk of the molecular material. The FIR emission in the 70~\\arcsec~LWS beam at the nucleus is dominated by emission from star formation rather than AGN activity. On scales of $\\sim$ 1 kpc the average physical properties of the PDRs are modelled with a gas density, n $\\sim$ 10$^3$ cm$^{-3}$, an incident far-UV field, G $\\sim$ 10$^2$ and a gas temperature of $\\sim$ 250 K."
    },
    "0002/astro-ph0002354_arXiv.txt": {
        "abstract": "We present the results of a systematic search for X-ray flares on young stars observed during {\\em ROSAT} PSPC observations of the Taurus-Auriga-Perseus sky region. All pointed PSPC observations currently available from the {\\em ROSAT} Public Data Archive with known pre-main sequence T Tauri Stars or young Pleiads or Hyads in the field of view are analyzed. A study of the activity of late-type stars of different ages provides information on the evolution of their coronal activity, which may be linked to their angular momentum.  We develop a criterion for the detection of flares based on the shape of the X-ray lightcurve. Applying our detection method to all 104 PSPC pointings from the archive we find 52 flares. Among them 15 are detected on T Tauri Stars, 20 on Pleiads, and 17 on Hyads. Only the 38 events which can definitely be attributed to late-type stars (i.e. stars of spectral type G and later) are considered in the statistical analysis of the properties of flaring stars. We investigate the influence of stellar parameters such as age, rotation and multiplicity on individual flare parameters and flare frequency.  From the total exposure time falling to the share of each sample and the duration of the individual flares we compute a flare rate. We take into account that the detection sensitivity for large X-ray flares depends on the S/N and hence on the stellar distance. The values we derive for the flare rates are $0.86 \\pm 0.16$\\% for T Tauri Stars, $0.67 \\pm 0.13$\\% for Pleiads and $0.32 \\pm 0.17$\\% for Hyads. The flare rate of classical T Tauri Stars may be somewhat higher than that of weak-line T Tauri Stars ($F_{\\rm c} = 1.09 \\pm 0.39$\\% versus $F_{\\rm w} = 0.65 \\pm 0.16$\\%).  Hardness ratios are used to track the heating that takes place during stellar flares. Hardness ratios are evaluated for three distinct phases of the flare: the rise, the decay, and the quiescent (pre- and post-flare) stage. In most cases the hardness increases during the flares as compared to the quiescent state. During both quiescence and flare phase TTSs display the largest hardness ratios, and the Hyades stars show the softest spectrum. ",
        "introduction": "\\label{sect:intro}  The Taurus-Auriga-Perseus region offers the opportunity to study the X-ray emission of young stars at several evolutionary stages. The youngest stars observed by {\\em ROSAT} in this portion of the sky are the T Tauri Stars (TTSs) of the Taurus-Auriga and Perseus star forming regions, late-type pre-main sequence (PMS) stars of $M \\leq 3 {\\rm M}_\\odot$ with an estimated age of $10^5-10^7\\,{\\rm yrs}$. Two young star clusters, the Pleiades and Hyades, are also located in this region of the sky at age of $10^8\\,{\\rm yrs}$ and $6~10^8\\,{\\rm yrs}$, respectively. They consist mostly of zero-age main-sequence (ZAMS) stars, except for some higher mass post-main sequence stars and brown dwarfs, which are not studied here.  From the early observations by the {\\em Einstein} satellite it was concluded that the X-ray emission of young stars arises in an optically thin, hot plasma at temperatures above $10^6\\,{\\rm K}$ (\\cite{Feigelson81.1}). The emission region has been associated with the stellar corona where the X-rays are produced --- more or less analogous to the solar X-ray emission --- through a stellar $\\alpha$-$\\Omega$-dynamo. The dynamo is driven by the combination of rotation and convective motions. Correlations between the X-ray emission of late-type stars and the stellar rotation support the notion that dynamo-generated magnetic fields are responsible for heating the coronae (\\cite{Pallavicini81.1}). But successful direct measurements of the magnetic fields of TTSs have been performed only recently (see e.g. \\cite{Guenther99.2}). The details of the heating mechanism are still not well understood.  The correlation between stellar rotation and X-ray emission of late-type stars suggests that the rotational evolution of young stars determines the development of stellar activity. The rotational evolution of low-mass PMS stars partly depends on the circumstellar environment. While classical TTSs (hereafter cTTSs) are surrounded by a circumstellar disk, inferred from IR dust emission (\\cite{Bertout88.1}, \\cite{Strom89.1}, and \\cite{Beckwith90.1}) and more recently from direct imaging (e.g. \\cite{McCaughrean96.1}), weak-line TTSs (wTTSs) lack such a disk, or at least the disk is not optically thick. Owing to contraction wTTSs spin up as they approach the main sequence. For cTTSs, on the other hand, coupling between the disk and the star may prevent spin-up (\\cite{Bouvier93.1}). The period observed on the ZAMS depends on the time the star has spent in the cTTS phase. After the main-sequence is reached, the rotation rate decreases again (see \\cite{Bouvier97.1}). As a consequence of their slower rotation, stars on the ZAMS and main sequence (MS) should on average show less X-ray activity than PMS stars.  Earlier investigations of X-ray observations of young late-type stars were mostly concerned with the quiescent emission (see \\cite{Neuhaeuser95.1}, \\cite{Stauffer94.1}, \\cite{Gagne95.1}, \\cite{Hodgkin95.1}, \\cite{Micela96.1}, 1999, \\cite{Pye94.1}, and \\cite{Stern94.1}). In contrast to these studies we focus on the occurrence of X-ray flares. Furthermore we discuss a larger sample than most of the previous studies by using {\\em all} currently available observations from the {\\em ROSAT} Public Data Archive that contain any TTS, Pleiad or Hyad in the field of view.  X-ray flares may be used as a diagnostics of stellar activity. They are thought to originate in magnetic loops. In contrast to findings from quasi-static loop modeling, the only direct determination of the size of a flaring region (\\cite{Schmitt99.1}) shows that the emitting region is very compact. In the loops which confine the coronal plasma magnetic reconnection suddenly frees large amounts of energy which is dissipated into heat and thus leads to a temporary enhancement of the X-ray emission. The decay of the lightcurve is accompanied by a corresponding (exponential) decay of the temperature and emission measure, which are obtained from one- or two-temperature spectral models for an optically thin, thermal plasma (\\cite{Raymond77.1}, \\cite{Mewe85.1}, 1986).  The most powerful X-ray flares have been observed on the youngest objects, notably a flare on the infrared Class I protostar YLW~15 in $\\rho$ Oph which has been presented by \\citey{Grosso97.1}. X-ray flares on TTSs observed so far (see \\cite{Montmerle83.1}, \\cite{Preibisch93.1}, \\cite{Strom94.1}, \\cite{Preibisch95.1}, \\cite{Gagne95.1}, \\cite{Skinner97.1}, \\cite{Tsuboi98.1}) exceed the maximum emission observed from solar flares by a factor of $10^3$ and more. Some extreme events have shown X-ray luminosities of $\\sim L_{\\rm x} = 10^{33}\\,{\\rm erg/s}$. Although some of the strongest X-ray flares ever observed were detected on TTSs to date no systematic search for TTS flares was undertaken.  This paper is devoted to a study of the relation between X-ray flare activity and other stellar parameters, such as age, rotation rate, and multiplicity. For this purpose we perform a statistical investigation of {\\em ROSAT} observations. We develop a method for the flare detection based on our conception of the typical shape of a flare lightcurve, where the term `typical shape' refers to the characteristics of the X-ray lightcurve described above, i.e. a significant rise and subsequent decay of the lightcurve to the previous emission level. The database and source detection is described in Sect.~\\ref{sect:data}. In Sect.~\\ref{sect:lcs} we describe how the lightcurves are generated. Our flare detection algorithm is explained in Sect.~\\ref{sect:detect}, where we also present all flare parameters derived from the X-ray lightcurves. Then we describe the influence of observational restrictions on the data analysis and how the related biases can be overcome (Sect.~\\ref{sect:bias}). In Sect.~\\ref{sect:statcomp} we compare the flare characteristics of different samples of flaring stars selected by their age, rotation rate, and multiplicity. We present luminosity functions for TTSs, Pleiads, and Hyads during flare and quiescence. Luminosity functions of the non-active state of these stars have been presented before (see e.g. \\cite{Pye94.1}, \\cite{Hodgkin95.1}, \\cite{Neuhaeuser97.3}) and some of the flares discussed here have been discussed by \\citey{Gagne95.1}, \\citey{Strom94.1}, and \\citey{Preibisch93.1}. However, this is the first statistical evaluation of flare luminosities. Flare rates comparing stellar subgroups with different properties (such as age, $v\\,\\sin{i}$, and stellar multiplicity) are compiled in Sect.~\\ref{sect:rate}. Because of lack of sufficient statistics for a detailed spectral analysis, hardness ratios are used to describe the spectral properties of the flares. In Sect.~\\ref{sect:hr} we present the observed relations between hardness ratios measured during different activity phases and between hardness and X-ray luminosity. Finally, we discuss and summarize our results in Sect.~\\ref{sect:discussion} and Sect.~\\ref{sect:conclusions}. ",
        "conclusions": "\\label{sect:conclusions}  We have determined flare rates for PMS stars, Pleiades and Hyades on a large data set and found that all stars are observed during flares for less than 1\\% of the observing time. Both frequency and strength of large X-ray flares decline after the PMS phase.  To probe whether the activity changes in the presence of a circumstellar disk, e.g. as a result of magnetic interactions between the star and the disk, we have compared flares on cTTSs and wTTSs. We find that flares on cTTSs are stronger and more frequent.  A comparison of flares on spectroscopic binaries to flares on all other stars of our sample shows that the flare rate is by a factor of $\\sim 2$ higher for the close binaries.  The flare rate of fast rotators is enhanced by a factor of $\\sim$ 3 as compared to slowly rotating stars.  To summarize, our analysis confirms that age and rotation influence the magnetic activity of late-type stars. All previous studies in this field have focused on the quiescent X-ray emission. Now, for the first time the rotation-activity-age connection has been examined for X-ray flares. Furthermore, from the sample of flares investigated here we find evidence that magnetic activity goes beyond solar-type coronal activity: On young stars interactions between the star and a circumstellar disk or the magnetic fields of close binary stars may play a role."
    },
    "0002/astro-ph0002448_arXiv.txt": {
        "abstract": "Supernovae in distant galaxies that are gravitationally lensed by foreground galaxy clusters make excellent cosmological candles for measuring quantities like the density of the Universe in its various components and the Hubble constant.  Distant supernovae will be more easily detectable since foreground cluster lenses would magnify such supernovae by up to 3--4 magnitudes. We show that in the case of the lens cluster Abell~2218, the detectability of high-redshift supernovae is significantly enhanced due to the lensing effects of the cluster. Since lensed supernovae will remain point images even when their host galaxies are stretched into arcs, the signal-to-noise ratio for their observation will be further enhanced, typically by an order of magnitude. We recommend monitoring well-modelled clusters with several known arclets for the detection of cosmologically useful SNe around $z=1$ and beyond. ",
        "introduction": "Observations of distant sources with known absolute luminosity (cosmological standard candles) are of primary importance to modern cosmology, since the relation between the apparent magnitude, luminosity and redshift of distant galaxies can be used to determine the Hubble constant $H_0$, the deceleration (or density) parameter $q_0$, and the cosmological constant $\\Lambda$.  Observations of standard candles beyond $z=0.05$ (where peculiar velocities are small) can yield the value of $H_0$ with reasonably small uncertainty (\\eg\\ \\cite{fil1,ham96}).  Observations of standard candles at $z\\!>\\!0.3$ are being used for the determination of the fraction of the total energy of the Universe in matter $\\Omega_M$ and in some hitherto unknown form $\\Omega_\\Lambda$ (\\cite{riess98,perlmutter99,saini00}).  The study of the gravitational magnification of standard candles at even higher redshift will put tighter constraints on dark matter models of cosmogony (\\eg\\ \\cite{kb98, marri98, holz98, metcalf99, pm00}).  The work of Riess \\etal\\ (1998b) and Perlmutter \\etal\\ (1999) has shown that, if detected significantly earlier than the epoch of their peak luminosity, type~Ia supernovae (SNe Ia) would be the most useful among cosmological candles at high redshift. However, the required integration times for good photometry and for obtaining spectra of such supernovae at redshift $z\\ssim 1$ are estimated to be tens of hours on a 10m telescope for $0\\farcs 75$ seeing (\\cite{goop95}). These observations would clearly be more favourable if these supernovae occur in galaxies magnified by gravitational lensing.  The magnification due to lensing can be significant enough to make possible the detection of supernovae (SNe) in galaxies at high redshifts ($z\\!\\gta\\!1$). Narasimha \\& Chitre (1988) first pointed out that such events in giant luminous arcs (as in the A370 system) can be used as a test of the lens models.  In the case of multiply imaged supernovae, Kovner \\& Paczy\\'nski (1988) deduce simple relations between the magnification of such a SN, the separation of images, and the differences between the arrival times of the event in different images.  Indeed, such SNe would serve as a unique probe for not only the distribution of matter in the clusters, but also for studying the source galaxies themselves. Due to the increased flux produced by the magnification of the images, photometric and spectroscopic studies of very distant galaxies can become possible.  This would enable us to obtain information, which would be otherwise unavailable, about the star formation process in the young galaxies (\\cite{mel91,yee}), the evolutionary status of AGN (\\cite{sti}), and even the morphology of distant galaxies (\\cite{ctt}). Indeed, one of the farthest known galaxies (at $z$=4.92, \\cite{fra97}) would not have been detected had it not been for the $10$-fold magnification by the cluster CL1358+62 at $z=0.33$.  In this paper we address the feasibility of detecting lensed SN events in high redshift galaxies which would be useful in the measurement of cosmological parameters. From a qualitative point of view such a study seems worthwhile for several reasons.  For a typical magnification of 3--4 mag (\\cite{kovpac88}) the study of lensed SNe stretches the usefulness of using them to characterize the distance ladder to further distances by a factor of 4--6, or, equivalently, results in a considerable decrease in the required duration of observation. Furthermore, although galaxies lensed into arcs are resolved in one direction due to stretching, a supernova in such a galaxy will remain a point source, hence the signal-to-noise ratio (\\snr) of a lensed supernova in an arc will be superior to that of one in an unlensed galaxy.  Finally, a cluster lens typically produces multiple images with time delays between them being up to several months, thereby making it possible to observe the same SN again, and measuring its light curve more accurately, particularly in its pre-peak phase.  In the searches we propose, we do not have to be confined to one lensed SN at a time. In many known cases of gravitational lensing of background objects by galaxy clusters, several arcs and arclets can be found in an area of the sky typically imaged by a single CCD frame.  In the case of \\object{Abell 2218}, for instance, there are 30 observed arclets (\\cite{ebb98, bez98}) with $R\\le 23.5$ and $\\mu_R\\le 25.5$ between $z=$0.5 to 1.5, so clearly a lot of galaxies can be simultaneously monitored. This is also true of the cluster \\object{Abell~2390}  at $z=0.23$, in which, in addition to the famous ``straight arc'' (triple image of a galaxy at $z=0.913$), there are at least 12 arclets ($R<21$) between $z=$0.4--1.3 in an area of $2.7\\times 2.7$ arcmin$^2$ around it (\\cite{bezecourt2390}). In the same area, in the magnitude range $21<R<23.5$, there are four images of two galaxies at $z=4.05$ (\\cite{pello2390,frye2390})  lensed by the same cluster.  Similarly, in a single HST WFPC field, which covers much less area ($\\sim$5 arcmin$^2$) than most CCD cameras on terrestrial telescopes, one finds 20 arclets brighter than $m_{\\rm F675W}$=24 (corresponding to at least 15 independent galaxies) beyond $z\\!=\\!0.7$ in the cluster \\object{Abell~370} ($z\\!=\\!0.37$, \\cite{bezecourt370}), and 14 arclets corresponding to 5 galaxies between $z=$0.5--1.5 in the cluster \\object{MS0440+0204} at $z\\!=\\!0.197$ (\\cite{gio98}). Not all of these, of course, would be magnified to the same degree, but on average they would be magnified, making them easier to be observed than in an unlensed case.  \\begin{figure}[t]   % \\centerline{ \\resizebox{\\hsize}{!}{\\rotatebox{-90}{\\includegraphics{snplot.ps}}}} \\caption[ ]{ The corrected apparent peak B~magnitudes of 79 published Type~Ia supernovae (Perlmutter \\etal\\ 1999, Riess \\etal\\ 1998, Hamuy \\etal\\ 1996) are plotted against redshift. The filled boxes represent SNe that have been used elsewhere (\\cite{perlmutter99,saini00}) to determine cosmological parameters, while the open boxes indicate SNe that are generally left out (including SN1997ck at $z=0.97$ which has not been spectroscopically confirmed to be type~Ia).  The lines represent models with ($\\Omega_M$, $\\Omega_\\Lambda$): (a) Standard Cosmology (1.0, 0.0), full line, (b) Open matter-only Universe (0.3, 0.0), dashed line (c) Perlmutter \\etal\\ best flat model: (0.28, 0.72), dot-dashed line, and (d) Perlmutter \\etal\\ best general model: (0.73, 1.32), dotted line. \\label{snplot}} \\end{figure}  Hereafter, we briefly review the usefulness of SN~Ia and SN~IIL in the determination of cosmological parameters, and in \\S3 quantify the effect of a cluster lens on the detectability of high-$z$ SNe.  We illustrate our case with the cluster lens Abell~2218.  In \\S4, we show that, in addition to this, for SNe detected in giant arcs, the signal-to-noise ratio for the detection of SNe is enhanced by an order of magnitude due to lensing. We summarize in \\S5 in the light of ongoing SN searches. ",
        "conclusions": "A considerable fraction of SNe in high redshift galaxies can be magnified by foreground galaxy clusters.  This pushes the equivalent observational distance to SNe further by a factor of 2--3, and thus allows measurements of magnitudes and light-curves of high redshift SNe in significantly shorter observational periods.  Even a small telescope with a limiting magnitude of $m_{\\rm lim}=24$ will detect SNe~Ia up to $z\\ssim 1.4$ in the field of a lens like A2218, while in the absence of such a lens, the same setup would not be able to detect SN~Ia beyond $z\\!=\\!0.7$ (Figure~2).  SNe of Type~II will never be detected by such a setup beyond $z\\!=\\!0.5$ for $m_{\\rm lim}<26$, while in the presence of a A2218-like lens, they could be detected up to $z\\ssim 2$.  SNe occurring in giant arcs will have an additional gain of signal-to-noise of $\\sim$15, making it easier to observe them.  The magnification for lensed SNe does not change dramatically for sources between $z=1\\!-\\!2$, while the surface brightness of the host galaxy decreases as $(1+z)^{-4}$, which in turn further improves the signal-to-noise ratio.  Such SNe will also be multiply imaged, further constraining mass models from time-delay measurements.  A number of projects have been searching for SNe at moderate and high redshifts. Neither of the high-$z$ SN search projects (\\cite{perlmutter99, schmidt98}) has yet discovered a SN with a favourable geometry behind a cluster. The low-$z$ Abell cluster search (\\cite{stromlo}) considers clusters out to $z=0.08$, again not ideal for these searches.  In order to find high-$z$ SNe without having to wait for the commissioning of NGST, one needs to continually monitor clusters with known arcs with a favourable geometry (\\eg\\ Abell 2218, Abell 370, CL0024+17) in the spirit of the SN Cosmology searches.  Once a SN event is observed, one needs to calculate the magnification (weakly model-dependent) from the location of the SN in the arc.  This will yield the intrinsic apparent magnitude of the SN from the light curve, which will enable us to calculate the distance to the supernova independent of its redshift.  For a SN in the redshift range $z=1\\!-\\!2$, this method can even yield a reliable value of $q_0$, because we would be measuring $H_0$ in the relativistic regime. As Fig.~{\\ref{snplot}} shows, in this redshift range, SNe~Ia with measured distances can distinguish, for example, between $\\Lambda=0$ and $\\Lambda\\ne0$ models even if the lens model does not allow the estimation of the magnification to better than 0.5 mag."
    },
    "0002/astro-ph0002162_arXiv.txt": {
        "abstract": "A peak has been unambiguously detected in the cosmic microwave background (CMB) angular spectrum.  Here we characterize its properties with fits to phenomenological models.  We find that the TOCO and BOOM/NA data determine the peak location to be in the range 175--243 and 151--259 respectively (both ranges 95\\% confidence) and determine the peak amplitude to be between $\\approx70~{\\rm and}~90$ $\\muK$.  By combining all the data, we constrain the full-width at half-maximum to be between 180 and 250 at 95\\% confidence.  Such a peak shape is consistent with inflation-inspired flat, cold dark matter plus cosmological constant models of structure formation with adiabatic, nearly scale-invariant initial conditions. It is inconsistent with open and defect models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002481_arXiv.txt": {
        "abstract": "We report cosmic-ray proton and helium spectra in energy ranges of 1 to 120~GeV and 1 to 54~GeV/nucleon, respectively, measured by a balloon flight of the BESS spectrometer in 1998. The magnetic-rigidity of the cosmic-rays was reliably determined by highly precise measurement of the circular track in a uniform solenoidal magnetic field of 1~Tesla. Those spectra were determined within overall uncertainties of $\\pm 5$~\\% for protons and $\\pm 10$~\\% for helium nuclei including statistical and systematic errors. ",
        "introduction": "The absolute fluxes and spectra of primary cosmic-ray protons and helium nuclei are fundamental information as references in cosmic-ray physics. They are needed to calculate secondary anti-protons, positrons, and diffuse gamma radiation, which in turn provide important knowledge of particle propagation and the matter distribution in interstellar space. Those are also indispensable for studying atmospheric neutrinos. Although measurement of the proton and helium energy spectra has been performed in various experiments, their resultant absolute fluxes show discrepancies up to a factor of two at 50~GeV. This ambiguity causes large uncertainty in calculations of the atmospheric neutrinos, as well as secondary anti-protons, positrons, and diffuse gamma-rays.  We report a new precise measurement of the cosmic-ray proton and helium spectra over the energy ranges of 1 to 120~GeV for protons and 1 to 54~GeV/nucleon for helium nuclei, based on a half of the data from a BESS balloon flight in 1998. The covered energy range is relevant to the atmospheric neutrinos observed as ``fully contained events'' in Super-Kamiokande. In the BESS--98 flight, a new trigger mode was implemented with a silica-aerogel Cherenkov counter to record all energetic particles instead of sampling the protons at a ratio of 1/60 as done in the previous BESS flights. This drastically improved statistics in the high-energy region above 6~GeV/nucleon, as reported here. ",
        "conclusions": "The proton and helium fluxes at the top of the atmosphere have been obtained from the BESS--98 flight data as summarized in Table~\\ref{tbl:result} and as shown in Figure~\\ref{fig:pftoa} in comparison with the previous experiments. The first and second errors in Table~\\ref{tbl:result} represent statistical and systematic errors, respectively. The overall errors including both errors are less than $\\pm 5$~\\% for protons and $\\pm 10$~\\% for helium nuclei. The dotted lines in Figure~\\ref{fig:pftoa} indicate the spectra assumed in the calculation of atmospheric neutrinos \\citep{hkkm}. Our results, as well as other recent measurements, are favorable to lower fluxes than the ones assumed in the atmospheric neutrino calculation especially above a few tens of GeV. It may suggest an importance of the reconsideration for the atmospheric neutrino flux predictions. Precise measurements of primary cosmic-rays will help to improve the accuracy in the atmospheric neutrino calculations."
    },
    "0002/hep-ex0002018_arXiv.txt": {
        "abstract": "I summarize the current state of solar neutrino research and then give my answer to the question: What should we do next? ",
        "introduction": "The reader who is familiar with solar neutrino research may wish to skip directly to the last section entitled: What Next?  Solar neutrinos have been detected experimentally with fluxes and energies that are qualitatively consistent with solar models that are constructed assuming that the sun shines by nuclear fusion reactions. The first experimental result, obtained by Ray Davis and his collaborators in 1968,~\\cite{davis68,chlorine} has now been confirmed by four other beautiful experiments, Kamiokande,~\\cite{kamiokande} SAGE,~\\cite{sage} GALLEX,~\\cite{gallex} and SuperKamiokande.~\\cite{superk} The observation of solar neutrinos with approximately the predicted energies and fluxes establishes empirically the theory~\\cite{bethe39} that main sequence stars derive their energy from nuclear fusion reactions in their interiors and has inaugurated what we all hope will be a flourishing field of observational neutrino astronomy.  Although the calculated neutrino fluxes depend upon high powers of the central temperature of the solar model, the experiments and the solar model theory are so precise that persistent quantitative discrepancies have existed between the model predictions and the solar model calculations for over thirty years~\\cite{bahcall68,bahcall89,BP98}  Important experiments are underway that will provide diagnostic information about the physical properties of neutrinos that are created in the center of the sun and detected on earth in really long baseline experiments.  At this workshop, we will hear discussions of the SuperKamiokande, SNO, BOREXINO, \\hbox{HELLAZ}, HERON, ICARUS, LENS, and KamLAND experiments.      I will discuss predictions of the combined standard model in the main part of this review.  By `combined' standard model, I mean the predictions of the standard solar model and the predictions of the standard electroweak model.  We need a solar model to tell us how many neutrinos of what energy are produced in the sun and we need an electroweak theory to tell us how the number and flavor content of the neutrinos are changed as they make their way from the center of the sun to detectors on earth.  For all practical purposes, the standard electroweak model states that nothing happens to solar neutrinos after they are created in the deep interior of the sun.  Using standard electroweak theory and fluxes from the standard solar model, one can calculate the rates of neutrino interactions in different terrestrial detectors with a variety of energy sensitivities.  The combined standard model also predicts that the energy spectrum from a given neutrino source should be the same for neutrinos produced in terrestrial laboratories and in the sun and that there should not be measurable time-dependences (other than the seasonal dependence caused by the earth's orbit around the sun).  The spectral and temporal departures from standard model expectations are expected to be small in all currently operating experiments~\\cite{bks98} and have not yet yielded definitive results. Therefore, I will concentrate here on inferences that can be drawn by comparing the total rates observed in solar neutrino experiments with the combined standard model predictions.  I will begin by reviewing in Section~\\ref{best} the quantitative predictions of the combined standard solar model and then describe in Section~\\ref{threeproblems} the three solar neutrino problems that are established by the chlorine, Kamiokande, SAGE, GALLEX, and SuperKamiokande experiments.  In Section~\\ref{uncertainties}, I detail the uncertainties in the standard model predictions and then show in Section~\\ref{helioseismology} that helioseismological measurements indicate that the standard solar model predictions are accurate for our purposes. In Section~\\ref{helioseismology}, I discuss the implications for solar neutrino research of the precise agreement between helioseismological measurements and the predictions of standard solar models. Next, ignoring all knowledge of the sun, I cite analyses  in Section~\\ref{nomodels} that show that one cannot fit the existing experimental data with neutrino fluxes that are  arbitrary parameters, unless one invokes new physics to change the shape or flavor content of the neutrino energy spectrum. I summarize in Section~\\ref{oscillations} the characteristics of the best-fitting neutrino oscillation descriptions of the experimental data. Finally, I will discuss and summarize the results in Section~\\ref{discussion}.  If you want to obtain  numerical data or subroutines that are discussed in this talk, or to  see relevant background information, you can copy  them from my Web site: http://www.sns.ias.edu/$\\sim$jnb.  Before we begin the detailed discussion, I want to make just a brief historical diversion.  Nearly all of the current interest in solar neutrinos centers around the opportunity to use the sun as a neutrino source in a very long baseline oscillation experiment.  In preparing for a talk in honor of Fred Reines a few months ago, I ran across a long forgotten 1972 letter from Bruno Pontecorvo, the originator of the hypothesis that oscillations may be observed in solar neutrino experiments. For your interest, I enclose a reproduction of this letter in Fig.~\\ref{fig:pontecorvo}.  \\begin{figure}[!t] \\centerline{\\psfig{figure=bahcallfig1.ps,width=3in}} \\vglue-.3in \\caption[]{Letter from Bruno Pontecorvo in 1972.\\label{fig:pontecorvo}} \\end{figure}   For the benefit of neutrino pioneers of today, it is perhaps worth remarking that Ray Davis and I never considered the possibility that solar neutrinos could be used to learn more about neutrinos when, in the early 1960's, we were first analyzing the potentialities of a practical chlorine experiment.  We sold the experiment as a fundamental  test of the hypothesis that the sun shines by nuclear fusion reactions in its interior.  Only after the first results of the chlorine experiment showed in 1968 a conflict with the solar model calculations and Gribov and Pontecorvo published their epochal 1969 paper on vacuum oscillations of solar neutrinos did we begin to consider the possibility that solar neutrinos might tell us something new about particle physics. Maybe there are previously unimagined physics treasures to be discovered in future neutrino experiments. ",
        "conclusions": "\\label{discussion}  When the chlorine solar neutrino experiment was first proposed,~\\cite{davis64} the only stated motivation was ``...to see into the interior of a star and thus verify directly the hypothesis of nuclear energy generation in stars.''  This goal has now been achieved,  The focus has shifted to using solar neutrino experiments as a tool for learning more about the fundamental characteristics of neutrinos as particles.  Experimental effort is now concentrated on answering the question: What are the probabilities for transforming a solar $\\nu_e$ of a definite energy into the other possible neutrino states? Once this question is answered, we can calculate what happens to $\\nu_e$'s that are created in the interior of the sun.  Armed with this information from weak interaction physics, we can return again to the original motivation of using neutrinos to make detailed, quantitative tests of nuclear fusion rates in the solar interior. Measurements of the flavor content of the dominant low energy neutrino sources, $p$-$p$ and $^7$Be neutrinos, will be crucial in this endeavor and will require another generation of superb solar neutrino experiments (see the comments in Section~\\ref{whatnext}).   Three decades of refining the input data and the solar model calculations has led to a predicted standard model event rate for the chlorine experiment, $7.7$ SNU, which is very close to $7.5$ SNU, the best-estimate value obtained in 1968.~\\cite{bahcall68} The situation regarding solar neutrinos is, however, completely different now, thirty years later. Four experiments have confirmed the original chlorine detection of solar neutrinos. Helioseismological measurements are in excellent agreement with the standard solar model predictions and very strongly disfavor (by a factor of $40$ or more) hypothetical deviations from the standard model that are require to fits the neutrino data (cf.~Fig.~\\ref{modelsunnu}). Just in the last two years, improvements in the helioseismological measurements have resulted in a five-fold improvement in the agreement between the calculated standard solar model sound speeds and the measured solar velocities (cf. Figure~2 of the Neutrino~96 talk~\\cite{neutrino96} with Figure~\\ref{modelsunnu} of this talk)."
    },
    "0002/astro-ph0002432_arXiv.txt": {
        "abstract": "We present a detailed study of the morphological features of 22 rich galaxy clusters. Our sample is constructed from a cross-correlation of optical (Abell+APM) data with X--ray (0.1 - 2.4) keV ROSAT pointed observations. We systematically compare cluster images and morphological parameters in an attempt to reliably identify possible substructure in both optical and the X--ray images. To this end, we compute various moments of the optical and X--ray surface-brightness distribution such as the ellipticities, center-of-mass shifts and ellipsoidal orientations. We assess the significance of our results using Monte Carlo simulations. We find significant correlations between the optical and X--ray morphological parameters, indicating that in both parts of the spectrum it is possible to identify correctly the dynamical state of a cluster. Most of our clusters (17/22) have a good 1-to-1 correspondence between the optical and the X--ray images and about 10 appear to have strong indications of substructure. This corresponds to a minimum percentage of order $\\sim 45 \\%$, which is in very good accordance with other similar analyses. Finally, 5 out of 22 systems ($\\sim 22 \\%$) seem to have distinct subclumps in the optical which are not verified in the X--ray images, and thus are suspect of being due to optical projection effects. These results will serve as a useful guide in interpreting subsequent analyses of large optical cluster catalogues.  {\\bf Keywords:} galaxies: clusters: general - large-scale structure of universe -  X--ray: galaxy clusters ",
        "introduction": "Galaxy clusters occupy a special position in the hierarchy of cosmic structure in many respects. Being the largest physical laboratories in the universe, they appear to be ideal tools for studying large-scale structure, testing theories of structure formation and extracting invaluable cosmological information, especially regarding the Hubble parameter and the value of $\\Omega_{\\circ}$ (cf. B${\\rm \\ddot{o}}$hringer 1995; West, Jones \\& Forman 1995; Buote 1998; Schindler 1999).  One of the most significant properties of galaxy clusters is the relation between their dynamical state and the underlying cosmology. In an open universe, clustering effectively freezes at high redshifts ($z\\simeq\\;\\Omega_{\\circ}^{-1}-1$) and clusters today should be more relaxed with weak or no indications of substructure. Instead, in a critical density model, such systems continue to form even today and should appear to be dynamically active. Even, in a small but well-controlled (from any kind of selection or other biases) cluster sample, the percentage and morphologies of disordered and perturbed objects could lead to constraints on the $\\Omega_{\\circ}$ and $\\Lambda$ parameters, especially if combined with N-body/gasdynamic numerical simulations spanning different dark matter (DM) scenarios (cf. Richstone, Loeb \\& Turner 1992 hereafter RLT92; Evrard et al. 1993; Lacey \\& Cole 1993).  The above pioneering works were the first to set some limits on $\\Omega_{\\circ}$ using the rate of cluster formation for various cosmologies and led the way to a plethora of research works towards this direction (see Figure 2 of RLT92). Since then, a large number of relevant analyses have been devoted to this study and an accordingly varying and large number of optical and X--ray cluster compilations have been utilised to this aim. All the methods employed in each case are quite different and all the related studies find an appreciable percentage of dynamically active galaxy clusters (see Forman \\& Jones 1990; B${\\rm \\ddot{o}}$hringer 1995; West 1995; Thomas et al. 1998 for good reviews of the subject). However, these studies do disagree on the precise number of clusters exhibiting significant dynamical activity, which varies between 30\\% and 80\\% of the total number of clusters studied, and which seems also to depend on the techniques employed in each analysis.  We quote only some of the most recent studies together with the sample used and the preference of the method given in each case. Geller \\& Beers (1982) analysed the iso-intensity contour maps of 65 optical clusters maps. Forman et al. (1981) studied 4 Einstein clusters. Dressler \\& Shectman (1988) studied velocity dispersion diagnostics for 15 optical clusters. Rhee, van Haarlem \\& Katgert (1991) and Solanes, Salvador-Sole \\& Gonzalez-Casado (1999) utilised 107 and 67 ENACS clusters respectively and a variety of 2D and 3D statistical tests to quantify the significance of substructure. Similarly, Dutta (1995), Crone, Evrard \\& Richstone (1996 hereafter CER96), Pinkney et al. (1996), Thomas et al. (1998 hereafter T98) and Jones \\& Forman (1999 hereafter JF99; 208 ROSAT X--ray clusters) proposed a variety of substructure tests depending on the level of information available (1D, 2D, 3D) and applied these to different N-body cluster data using several DM models. CER96 and T98 have explicitly argued that variations in the cluster center-of-mass as a function of distance or density (overdensity) threshold is one of the best possible substructure measures (see also West \\& Bothun 1990). The latter has previously been adopted by Mohr, Fabricant \\& Geller (1993; 5 X--ray clusters), Mohr et al. (1995; 65 Einstein clusters) and Rizza et al. (1998; 11 ROSAT HRI distant clusters). Evrard et al. (1993) and Mohr et al. (1993; 1995) were the first to make extensive use of the surface-brightness moments of the X--ray cluster distribution (eg. center-of-mass shifts). Gomez et al. (1997; 9 Abell clusters) have also employed the same methods and claimed that variations in the cluster ellipticity and orientation as a function of distance from the cluster center (isophotal twisting) should be considered as one of the prime substructure diagnostics. Bird (1994) in her sample of 25 Abell cluster has utilised galaxy peculiar velocities to locate and identify subclumps in the galaxy distribution. Kriessler \\& Beers (1997; 56 Abell and other clusters) have searched for cluster substructure signatures using the innovative KMM algorithm on the surface-density galaxy maps and quantified their findings using N-body simulations. Furthermore, Buote \\& Tsai (1995; 1996 hereafter BT95; BT96) and Valdarnini, Ghizzardi \\& Bonometto (1999) used the 2D gravitational potential moments (power ratio method) to characterise the dynamical state of clusters. Serna \\& Gerbal (1996; 2 Abell clusters) have developed the so-called hierarchical method to define and identify substructure in Abell clusters, while Slezak et al. (1994; 11 X--ray clusters) and Lazzati et al. (1998; 2 X--ray clusters) have adopted wavelet transform techniques in similar attempts.  As it is evident from all the above there is neither agreement on the methods utilised nor on the exact frequency of perturbed clusters. It seems that identifying significant dynamical activity within galaxy clusters, in close relation to the underlying cosmology and the density parameter, still remains an open issue. Not only do we need a large, statistically complete, cluster sample but an objective and reliable definition of what substructure is, upon which our study should be based.  The large majority of the analyses carried out so far, have made use of either optical (Abell) or X--ray (ROSAT, Einstein) cluster data. However, Mohr et al. (1993) and Rizza et al. (1998) have investigated cluster substructure using in a complementary fashion optical and X--ray data. In the present work, we extend this approach using a sample of 22 galaxy clusters for which we have data both in the optical and the X--ray part of the spectrum (APM and ROSAT respectively). Our sample size is twice as large as that of Rizza et al. (1998) and more than four times larger than that of Mohr et al. (1993). The advantage of using X--ray data is that the X--ray emission is proportional to the square of the gas density (rather than just density in the optical) and emanates mostly from the central cluster region, a fact which minimises projection effects (cf. Sarazin 1988; Schindler 1998). The advantage of using optical data is the shear size of the available cluster catalogues and thus the statistical significance of the emanating results. Subsequently, we are not only interested in comparing optical to X--ray cluster data regarding the various substructure tools, but to calculate and calibrate different biases using the superior X--ray data with which we could measure and test easily similar substructure diagnostics in a large, solely optical dataset as well. Therefore it is of great importance to address the following two questions, which we attempt to do in our present study:  \\begin{itemize} \\item Is substructure in the optical also corroborated by the X--ray observations and in what percentage? \\item  What is the confirmed percentage of systems depicting significant indications of internal activity? \\end{itemize}  To this aim we compare optical to X--ray morphological cluster parameters (position angles, ellipticities, centroid shifts and group statistics) in an attempt to classify objects according to their dynamical state and we compute the relative frequency of substructure.  We proceed by presenting the optical and X--ray datasets in section 2. In section 3, we describe our methods of analysis, in section 4 we present a comparison of the optical and X--ray cluster images while in section 5 we present the results of our substructure analysis. Finally, we draw our conclusions in section 6. Also, there is an extensive appendix at the end of this paper, where morphological and dynamical information for each individual cluster can be found as well as a comparison with other similar works on the common clusters. ",
        "conclusions": "We have investigated a sample of 22 galaxy clusters using optical (APM) and X--ray (ROSAT) data with the aim of addressing two questions: (a) Do optical and X--ray data reveal the same cluster morphological features ? (b) What is the percentage of relaxed and dynamically active clusters ?  Our cluster sample has flown from a cross-correlation between the APM, ACO and ROSAT pointed observations data and has a depth distribution with $z \\mincir 0.13$. We have examined our cluster sample utilising several cluster-morphology diagnostics such as isodensity maps, orientations, ellipticities, centroid variations and subgroup statistics. Looking at the isodensity contour maps, the cluster orientations and ellipticities, we observe a remarkable 1-to-1 correspondence between X--ray and optical data in $\\sim 80\\%$ of our sample, regarding the gross cluster characteristics (prime structures, elongations, multimodality, collision signatures). We quantify this by correlating the optical and X--ray cluster ellipticities and orientations and find high and statistically significant correlations. In an attempt to quantify the compatibility of the different substructure measures, that we have used, we correlate $\\epsilon$'s and $sc$'s, in the optical and in the X--ray separately and we also cross-correlate them to find significant and strong correlations. This implies that indeed the flattening of clusters is due to their dynamical activity.  From our substructure analysis we find that $\\sim 10$ out of 22 systems ($\\sim 45\\%$) display strong substructure indications visible in both parts of the spectrum. We also find that 5 clusters ($\\sim 22\\%$) show clear disparities between the optical and X--ray maps, with apparent substructure in the optical not corroborated by the available X--ray data. This is most possibly due to optical projection effects. Our results on the frequency of disordered clusters do concur with most of the relevant studies published to-date.  On the understanding that our catalogue is rather inadequate for drawing significant cosmological conclusions, we would prefer to be cautious when it comes to such a precarious task. Nevertheless, we observe that our present analysis is compatible with that of RLT92 (see their Figure 2) regarding the cluster substructure frequency, setting a rather frail lower limit on the density parameter ($\\Omega_{\\circ}\\,\\geq\\,0.5$; see also West 1995; West et al. 1995). Furthermore, the relatively large fraction (4/22) of recent mergers that we have identified is again indicative of a high-$\\Omega_{\\circ}$ Universe (see discussion of Zabludoff \\& Zaritsky 1995).  In the near future we plan to apply the methodology of this work to $\\sim 900$ APM galaxy clusters, in order to investigate in more detail the issue of cluster substructure."
    },
    "0002/astro-ph0002118_arXiv.txt": {
        "abstract": "Infrared $\\rm JK_{\\rm s}$-band photometry of the Galactic globular clusters $\\rm M\\,15$ and, for the first time, $\\rm M\\,56$ is presented. We estimated the reddening ($\\rm E(B-V)=0.18\\pm0.08$ mag) and distance modulus ($\\rm ({\\it m}-M)_V=15.43\\pm0.30$ mag) towards the poorly studied globular cluster $\\rm M\\,56$. We combined our data with observations of other clusters from the literature (12 in total) to extend the ${\\rm [Fe/H]}$ vs. Red Giant Branch (RGB) slope relation towards metal-poor clusters. Our best fit yields to ${\\rm [Fe/H]}\\,=\\,-3.40(\\pm0.22)\\,-27.74(\\pm2.35)\\,\\times\\,(RGB\\,Slope)$, with an $\\rm r.m.s.\\,=\\,0.20$. The broader metallicity baseline greatly reduced the uncertainties compared to other existing calibrations. We confirmed a previously obtained calibration of the relation between the RGB color $\\rm (J-K_{\\rm s})_0(RGB)$ at $\\rm M_{K_{\\rm s}}=-5.5$ vs. $\\rm[ Fe/H]$: ${\\rm [Fe/H]}\\,=\\,-6.90(\\pm0.99)\\,+6.63(\\pm1.05)\\,\\times\\,(J-K_{\\rm s})_0(RGB)$ with an $\\rm r.m.s.\\,=\\,0.33$. Finally, using the new RGB slope calibration we estimated the abundance of the super metal-rich cluster Liller 1 ${\\rm [Fe/H]}\\,=\\,0.34\\pm0.22$. ",
        "introduction": "Globular clusters are a fundamental laboratory for the study and understanding of stars and their evolution. They offer a unique opportunity to study samples of stars with a single age and metallicity, and provide a sequence of parameters which describe the stellar populations as ensembles of stars. Those parameters include the position of the red giant branch (RGB), the relative populations of the blue and red horizontal branches, and the relative number of RR Lyr variables.  Davidge et al. (1992\\markcite{dav92}) considered for the first time the slope of the RGB on a combined optical-infrared color-magnitude diagram (CMD) as a metallicity indicator. This technique is particularly promising (Kuchinski, Frogel \\& Trendrup 1995\\markcite{kuc95}; Kuchinski \\& Frogel 1995\\markcite{kuc95b}; Ferraro et al., 2000\\markcite{fer00}), because it is not affected by the reddening towards the cluster. Bypassing the reddening correction can help to improve significantly the photometric determination of the metallicity of more distant and highly obscured systems.  The stars on the RGB radiate most of their energy in the infrared, with the added advantage that the reddening is greatly diminished in this part of the spectrum. Although both those arguments are not relevant for observations of Galactic globular clusters, they become critical in the studies of distant systems, where the patchy internal extinction is added to the foreground extinction in the Milky Way, which makes the interpretation of the data more difficult. Also, many potential targets of interest (such as nearby dwarf galaxies) are expected to have a low metal content. Until recently, the available calibrations did not span all the necessary range of $\\rm[Fe/H]$. Indeed, Kuchinski \\& Frogel (1995\\markcite{kuc95b}) based their calibration on a set of Galactic globular clusters with metallicities from $\\rm[Fe/H]$~=~1.01 to -0.25. Ferraro et al. (2000\\markcite{fer00}) presented for the first time high quality near infrared CMDs of 10 Galactic globular clusters, and a detailed analysis of the RGB behavior as a function of metallicity.  The goal of this work is to increase the statistical basis of the RGB slope vs. $\\rm[Fe/H]$ calibration, and to test the existing calibrations (Kuchinski, Frogel \\& Trendrup 1995\\markcite{kuc95}; Kuchinski \\& Frogel 1995\\markcite{kuc95b}; Ferraro et al., 2000\\markcite{fer00}). Clusters with low metallicities are of particular interest, because they will allow to make this tool applicable to metal-poor stellar systems. We will also increase the statistical basis of the RGB slope vs. $\\rm[Fe/H]$ calibration.  We present here $\\rm JK_s$ photometry of the central area of $\\rm M\\,15$ and, for the first time, $\\rm M\\,56$. The basic data for the clusters are given in Table~\\ref{tbl-1} (Harris 1996\\markcite{har96}; June 22, 1999 version). They are both extremely metal-poor clusters. The table contains data on two more clusters ($\\rm M\\,30$ and $\\rm NGC\\,6553$) which we collected from the literature.  \\placetable{tbl-1}  $\\rm M\\,15$ is a well studied Galactic globular cluster. It possesses one of the highest known central densities (Yanny et al. 1994\\markcite{yan94}). King (1975\\markcite{kin75}) and Bahcall \\& Ostriker (1975\\markcite{bah75}) speculated that the cluster might have undergone a core collapse or might contain a central black hole. Sandage (1970\\markcite{san70}) obtained ground-based photometry of the outer region of $\\rm M\\,51$ and detected an extended blue horizontal branch (hereafter HB) but no significant population of blue stragglers. Subsequent photometric works were presented  by Auri\\`{e}re \\& Cordoni (1981\\markcite{aur81}), Buonano et al. (1985\\markcite{buo85}), Bailyn et al. (1988\\markcite{bai88}), and Cederbloom et al. (1992\\markcite{ced92}). More recently, the cluster was observed in the optical with the {\\it HST} by Ferraro \\& Parsce (1993\\markcite{fer93}) and Yanny et al. (1994\\markcite{yan94}). Frogel, Persson \\& Cohen (1983\\markcite{fro83}) published $\\rm JHK$ measurements of five bright red giants in $\\rm M\\,15$. The most recent variability study of $\\rm M\\,15$ (Buter et al., 1998\\markcite{but98}) reported light curves of 30 confirmed variable stars, mostly RR Lyr.  In contrast, $\\rm M\\,56$ is surprisingly poorly studied, probably because is lays close to the Galactic plane $(l=62.66\\arcdeg, b=+8.34\\arcdeg)$. Rosino (\\markcite{ros51}1951) obtained the first photographic CMD of this cluster. Barbon (1965\\markcite{bar65}) built the first CMD in the standard $BV$ colors. Smriglio, Dasgupta \\& Boyle (1995\\markcite{smr95}) used the Vilnius photometric system to estimate the extinction towards $\\rm M\\,56$, and pointed to a possible reddening variation across the cluster area. A number of observations of variable stars in $\\rm M\\,56$ have been undertaken throughout the years (Sawyer 1940\\markcite{saw40}; Sawyer 1949\\markcite{saw49}; Rosino 1961\\markcite{ros61}; Wehlau \\& Sawyer Hogg 1985\\markcite{weh85}). The latest CMD for this cluster (Grundahl et al. 1999\\markcite{gru99}) is in the Str\\\"{o}mgren $(u,u-y)$ system and shows very well defined blue and red HBs. ",
        "conclusions": "We showed that the position of the infrared RGB can be used to reliably determine the abundance of metal-poor stellar systems with an accuracy of $\\approx0.2$ (in $\\rm [Fe/H]$). Our main results include:  (1) We presented infrared photometry of the Galactic globular clusters $\\rm M\\,15$ and, for the first time, $\\rm M\\,56$, and studied the morphology of their CMDs.  (2) We estimated the reddening towards $\\rm M\\,56$ by comparing the RGB color at the level of HB to that of $\\rm M\\,15$, and obtained $\\rm E(B-V)=0.18\\pm0.08$ mag. We used the relative HB levels of the same two clusters to derive a distance modulus to $\\rm M\\,56$ of $\\rm ({\\it m}-M)_V=15.43\\pm0.30$ mag if $\\rm ({\\it m}-M)_V=15.37$ mag is assumed for $\\rm M\\,15$.  (3) We compiled a sample of 12 Galactic globular clusters with high quality infrared photometry. We recalibrated and extended the RGB slope vs. $\\rm [Fe/H]$ relation towards low metallicity globular clusters. We also reevaluated the RGB color $\\rm (J-K_{\\rm s})_0(RGB)$ at $\\rm M_{K_{\\rm s}}=-5.5$ mag vs. $\\rm[ Fe/H]$ relation. These are potentially useful tools to study extragalactic metal-poor stellar systems, particularly with high obscuration. Our results independently confirm the previously obtained calibrations.  (3) As an application, we used our newly determined RGB slope vs. $\\rm [Fe/H]$ relation to estimate the abundance of the super metal-rich Galactic globular cluster Liller 1, and obtained ${\\rm [Fe/H]}\\,=\\,0.34\\pm0.22$."
    },
    "0002/astro-ph0002097_arXiv.txt": {
        "abstract": "We present results from $\\lambda$ 3.6~cm observations of the quasar pair \\object{1038+528} A and B, made in 1995 using the VLBA together with the Effelsberg 100m telescope. We describe the use of a phase-referencing technique to measure the astrometric separation between the quasars. We also introduce a new data analysis method - \"hybrid double mapping\" - which preserves the relative astrometric information in a single VLBI hybrid map for close source pairs. We combine our measurements with those from three previous epochs, the earliest in 1981. Our new observations confirm the evolution within the structure of quasar B, previously proposed to explain the measured change in the relative separation of the pair. Our upper bound for any systematic proper motion between the mass centres of quasars A and B is 10~$\\mu$as~yr$^{-1}$. This is set by the limited precision in defining the reference points in the quasars at different epochs and by possible instabilities of the source \"core\" locations. A separate analysis enables us to put more stringent upper limits to any core motions along the two source axes. ",
        "introduction": "The quasar pair 1038+528~A,B (Owen et al. \\cite{owen78}) consists of two flat-spectrum radio sources, with redshifts 0.678 and 2.296 (Owen et al. \\cite{owen80}), separated on the sky by only 33\\arcsec. This system provides a unique opportunity to carry out high precision, relative astrometric studies using the full precision of VLBI relative phase measurements, since most sources of phase errors are common for the 2 sources (Marcaide \\& Shapiro \\cite{marca83}).  VLBI studies of the mas-scale structure of flat-spectrum quasars show that they typically have ``core-jet'' morphologies, consisting of a highly compact feature (the ``core'') located at the base of an extended linear feature or line of lower brightness components (the ``jet''). Both 1038+528 A and B exhibit such structures. In standard models of extragalactic radio sources, these radio-emitting features arise from a collimated beam of plasma which is ejected with a highly relativistic bulk velocity from a region close to a central massive object such as a black hole (see eg. Blandford \\& K\\\"onigl \\cite{blandford86}). Whilst jet features may correspond to shocks in the moving plasma, and can give rise to the observed ``superluminal'' component motions in some sources (Porcas \\cite{porcas87}), the ``core'' emission is thought to arise from a more-or-less permanent location close to the origin of the beam, where the ambient conditions correspond to a transition from optically thick to optically thin emission at the observed frequency. Although the ``core'' position may thus be frequency-dependent, for a fixed observing frequency the core should provide a stable marker, anchored to the central mass of the quasar, whose location can be used to define a precise position for the object as a whole. Although short time-scale variations in physical conditions may cause small changes in the ``core'' location, over long time-scales it may be used to track any systematic proper motion of the quasar.  The results from a near decade-long VLBI monitoring program on 1038+52A,B at $\\lambda$ 3.6 and 13~cm (from 1981.2 to 1990.5) are reported by Rioja et al. (\\cite{rioja97}), whose main conclusions can be summarized as follows:  \\begin{enumerate}  \\item There is no evidence of any relative proper motion between the quasars A and B. The uncertainties in the astrometric parameters result in an upper bound to any systematic relative motion between the cores of 10~$\\mu$as~yr$^{-1}$, consistent with zero.  \\item A compact feature within the jet of quasar B, chosen as the reference point for the structure, expands away from the core at a steady, slow rate of $\\sim~18~\\pm~5~\\mu$as~yr$^{-1}$, corresponding to v=$(0.8~\\pm~0.2)~h^{-1}$~c for a Hubble constant, H$_0$ = 100~h~km~s$^{-1}$~Mpc$^{-1}$; q$_0=0.5$. These values are used here throughout.  \\item The accuracy of the relative separation measurement is limited by noise and source structure, with estimated precisions of about $50~\\mu$as at $\\lambda$ 3.6~cm at any epoch.  \\item Confirmation of the consistently large offset (about 0.7 mas) between the positions of the peak of brightness (``core'') at $\\lambda$ 3.6 and 13~cm in quasar A.  \\end{enumerate}  New VLBI observations of this pair were made in November 1995 (1995.9) at $\\lambda$ 2, 3.6 and 13~cm. In this paper we report on results from our analysis of the 3.6~cm observations and investigate the temporal evolution of the source structures and relative separation from all four epochs spanning $\\sim 15$ years. Investigations of frequency-dependent source structure have also been made from a comparison of the astrometric measurements of the separations between A and B at all 3 wavelengths observed in 1995; these will be presented elsewhere (Rioja \\& Porcas in preparation).  Our new observations are described in Sect. 2. In Sect. 3 we describe the data reduction and mapping techniques used, and in Sect. 4 an analysis of the measurements in the maps. In Sect. 5 we compare the astrometric results from these observations with those from previous epochs and analyse the changes in separation. Conclusions are presented in Sect. 6. ",
        "conclusions": "The series of astrometric VLBI measurements of the separation between the quasar pair 1038+528 A and B, spanning nearly 15 years, provides excellent material for investigating the relative proper motion of two extragalactic radio sources and the positional stability of their cores. The changes measured in the separations between quasars A and B at 3.6~cm are dominated by the motion of the reference feature in quasar B. These astrometric results, and measurements in the hybrid maps of B are compatible with an expansion rate for the B reference component of 13--17$~\\mu$as~yr$^{-1}$. At a redshift of 2.296 this translates to an apparent transverse velocity of 0.55--0.70~c~h$^{-1}$. We note that this is an order of magnitude smaller than the more typical superluminal velocities seen in many quasars; it is a rare example of a subluminal velocity measured for a knot in a quasar jet.  After correcting for the motion of the reference component in B, we can put a conservative upper bound to any relative proper motion between the quasars of $10~\\mu$as~yr$^{-1}$. Despite the increase in temporal baseline, this upper bound is no better than that given by Rioja et al. (\\cite{rioja97}). Its value is related to the difficulty in reproducing a stable reference position along the A source axis near its ``core''.  Theories in which the redshifts of quasars do not indicate cosmological distances, and in which quasars are ``local'' and have high Doppler redshifts (e.g. Narlikar \\& Subramanian \\cite{narlikar83}) are incompatible with our measured upper-limit to relative proper motion. Quasars at 100 Mpc distance moving at relativistic speeds would have proper motions of the order of $600~\\mu$as~yr$^{-1}$, nearly 2 orders of magnitude greater than our limit. Assuming cosmological distances, our limit corresponds to apparent transverse velocities of 0.43~c~h$^{-1}$ and 0.22~c~h$^{-1}$ at the redshifts of the B and A quasars, respectively.  We have investigated the way in which the definition of reference points in a map may be only loosely ``fixed'' to the radio source structure, especially when the latter is strongly asymmetric. We have also developed an alternative analysis route - Hybrid Double Mapping - for imaging both sources of a close pair simultaneously, and at the same time preserving their relative astrometric information in a single map.  The surprisingly low rms from the fits of linear expansion in quasar B, and the discrepancy between the two estimates ($16.9~\\pm~0.6~\\mu$as~yr$^{-1}$ from the astrometric measurements and $ 13.0~\\pm~0.7~\\mu$as~yr$^{-1}$ from the hybrid maps) are suggestive of (but do not prove) a residual motion of the core in quasar B; our decomposition along the source axis direction gives a fit of $3.8~\\pm~0.3~\\mu$as~yr$^{-1}$, corresponding to an apparent transverse velocity of 0.17~c~h$^{-1}$. If real, this might indicate a steady change in physical conditions at the base of the jet, or perhaps the emergence of a new knot component moving outwards with a velocity similar to the reference component, but as yet unresolved by our 0.5 mas beam. In this regard, it is interesting to note the slight extension of the core of quasar B in PA 132\\degr given by the Gaussian model fit.  The low rms derived from fits to the expansion of the B reference component indicate that we have been overly conservative in our estimate of $18~\\mu$as for the error in reference point positions. Errors at least 2 times smaller are implied, corresponding to a sixtieth of the beamwidth. It is interesting to note that such small errors are also implied in the work of Owsianik \\& Conway (\\cite{owsianik98}), where the low scatter in the plot of expansion of the CSO source \\object{0710+439} allows an expansion rate of $14.1~\\pm~1.6~\\mu$as~yr$^{-1}$ to be determined.  There are no obvious systematic motions within quasar A, but the ``noise'' in the estimates of position along its axis are much larger. This noise, along with any associated underlying changes in source substructure, provides a fundamental limit to estimates of any systematic core motion in A. Improvements on the estimates of (or upper bounds to) the relative motion between the quasars, or of the individual motion of the A core, will require a considerable increase in the temporal baseline of VLBI monitoring.  \\newpage  \\appendix"
    },
    "0002/astro-ph0002262_arXiv.txt": {
        "abstract": "In this paper we define an observationally robust, multi-parameter space for the classification of nearby and distant galaxies. The parameters include luminosity, color, and the image-structure parameters: size, image concentration, asymmetry, and surface brightness. Based on an initial calibration of this parameter space using the ``normal'' Hubble-types surveyed by Frei \\etal (1996), we find that only a subset of the parameters provide useful classification boundaries for this sample.  Interestingly, this subset does not include distance-dependent scale parameters, such as size or luminosity. The essential ingredient is the combination of a spectral index (e.g., color) with parameters of image structure and scale: concentration, asymmetry, and surface-brightness.  We refer to the image structure parameters (concentration and asymmetry) as indices of ``form.'' We define a preliminary classification based on spectral index, form, and surface-brightness (a scale) that successfully separates normal galaxies into three classes. We intentionally identify these classes with the familiar labels of Early, Intermediate, and Late. This classification, or others based on the above four parameters can be used reliably to define comparable samples over a broad range in redshift. The size and luminosity distribution of such samples will not be biased by this selection process except through astrophysical correlations between spectral index, form, and surface-brightness. ",
        "introduction": "It is now well-established that a large fraction of galaxies discovered at intermediate and high redshift have unusual morphologies, and thus cannot be classified in terms of the nominal Hubble-Sandage system (Driver \\etal 1995, 1998; Abraham \\etal 1996a, 1996b). The Hubble classification scheme is also difficult to apply to many local galaxies, dubbed `peculiar,' or any galaxies imaged at low signal-to-noise ($S/N$), or apparently small size (relative to the point-spread function). The Hubble-Sandage classification system was predicated on the study of nearby, ``normal'' galaxies -- luminous and relatively quiescent objects(\\cite{san61}, \\cite{san87}, \\cite{san93}). While the classification system developed by de~Vaucouleurs \\etal (1976) makes an attempt to push the framework to ``later'' types, it still suffers from the above shortcomings. Fundamentally, these traditional classification schemes are based on the concept of `pigeon-holing' galaxies based on a reference-set, or archetypes. These archetypes are selected from samples in the local universe, and are preferentially axisymmetric systems.  Since our local census is undoubtedly incomplete, and, since galaxies evolve, such reference sets by their very definition are incomplete. Thus it is not surprising that these systems are of marginal utility in the study of dwarf galaxies, interacting galaxies, or galaxies at high redshift.  An alternative classification scheme could be based on quantitative indices, the inter-relation of which is not predetermined by a finite reference set. This would permit galaxies to be classified, for example, in different stages of their evolution; albeit the classification would be different but the basis set of indices would be the same. The goal of this paper is to define such a set of indices that can be used as quantitative, objective classifiers of galaxies (i) over a wide range in redshift, and (ii) for wide range of galaxy types. In particular, we desire classifiers that are well suited to typing both ``normal'' galaxies and the compact galaxies that are the focus of a companion study (Jangren \\etal 2000; hereafter, paper II). We anticipate that such a classification scheme is both necessary and enabling for the exploration of the physical mechanisms driving galaxy evolution (Bershady 1999).  What are the desired characteristics of classification parameters? They should be physically interesting (closely related to underlying physical properties of galaxies), model-independent, and measurable for all galaxy types. It also should be possible to accurately determine the parameters chosen for a wide range of image resolution and signal-to-noise ratios.  From Hubble's classification {\\it a posteriori} we have learned that a strong correlation exists between galaxy spectral type and apparent morphological features -- at least for the galaxy types which fit well within his scheme. This correlation -- noted by Hubble as early as 1936 (Hubble 1936) -- can loosely be termed a `color-morphology' relation, although the correlation is not necessarily limited to broad-band color. This is a triumph of Hubble's classification explicitly because it is not part {\\it of} the classification. Furthermore, the correlation yields clues about the physical connection of the present matter distribution and the star-formation histories in galaxies. But while morphology (or `form') and spectral type are correlated, there is also significant dispersion in this correlation. Some of the more notable deviations from the nominal color-morphology relation are found in the plethora of forms for spectrally `late' type galaxies, the presence of `E+A' galaxies (Dressler \\& Gunn 1993), and the compact, luminous, blue, emission-line galaxies studied in paper II (as we shall show). This points to the importance of form and spectral type as key, yet independent axes of a revised classification system.  However, the only example of such a revised classification system is that of Morgan (1958, 1959), where central light concentration is used as the primary classification parameter. Morgan was motivated by the fact that (i) a salient criterion used in classifying galaxies in the Hubble-Sandage system is the degree of central concentration of light; (ii) there was a significant dispersion in spectral type and Hubble type (\\cite{hum56}); and (iii) spectral type appeared to correlate more strongly with light concentration. In this way, Morgan hoped to wed the classification of stellar populations to the classification of galaxies. Nonetheless, he was compelled to introduce a secondary parameter, i.e. the `Form Family,' because there was still a dispersion of morphological forms within each of his spectral types. Today, one should be able to improve upon Morgan's scheme by introducing quantitative measures of image concentration and other indices of form, and by independently assessing the spectral type via colors or spectra.  A number of subsequent attempts have been made to construct quantitative classification system that could replace or modify the current Hubble scheme. Yet these schemes are generally based purely either on photometric form (e.g. \\cite{elm82}; \\cite{oka84}; \\cite{wat85}; \\cite{doi93}; \\cite{abr94}; \\cite{ode95}; \\cite{han95}) or spectral type (e.g. \\cite{ber95}; \\cite{conn95}; \\cite{zar95}; \\cite{fol96}; \\cite{bro98}; \\cite{ron99}). In essence, they have relied implicitly on an assumed correlation between galaxy spectral type and apparent morphology. Related attempts have been made to use artificial neural networks to reproduce the Hubble scheme in an objective way (e.g. \\cite{bur92}; \\cite{sto92}; \\cite{spie92}; \\cite{ser93}; \\cite{nai95}; \\cite{ode95}; \\cite{ode96}). Yet these go no further in differentiating between spectral type and form. Only in Whitmore's (1984) scheme are spectral and structural parameters combined, i.e., $B-H$ color, size, and bulge-to-total ratio are used to define two principal classification axes of {\\it scale} and {\\it form}. But again, the correlation(s) between galaxy spectral type, scale and form are not explicit.  Here we attempt to expand on Morgan's program by fully quantifying the classification of form via image concentration and several other structural parameters, and explicitly using color as an indicator of spectral type. In this study we choose to use only a single color ($B-V$), but we anticipate that a more desirable, future development would be to include broad-wavelength coverage, multi-color data and spectroscopic line-indices. Spectroscopic line-indices would be required, for example, to identify E+A galaxies. While such galaxies are not the focus of the present work, a comprehensive classification scheme should be able to isolate these systems and determine the range of their morphology ({\\it cf.} Dressler \\& Gunn 1992, Couch \\etal 1994, and Wirth \\etal 1994). Nonetheless, broad-band colors are a cost-effective way to characterize the spectral continuum ({\\it cf.} Bershady, 1995, and Connolly \\etal 1995). Of more direct relevance to the study at hand, a future elaboration of including $U-V$ and $V-K$ would enhance the ability to distinguish between spectral types particularly for galaxies with extremely blue, optical colors (e.g. Aaronson 1978, Bershady 1995).  We have also chosen to quantify form and scale via non-parametric measures, such as luminosity, half-light size and surface-brightness, asymmetry, and image concentration. An alternative, model-dependent approach is to decompose a galaxy's light profile into a disk and bulge. The traditional one-dimensional decompositions are fraught with technical problems such that decompositions can only be achieved reliably for about half of all disk galaxies (\\cite{ken85}). The newer two-dimensional decomposition techniques are superior (e.g. \\cite{jong96b}), and have been shown to successfully reproduce observed light profiles for faint galaxies (e.g. \\cite{sim99}). Indeed, one can argue that two-dimensional model fitting to imaging data is optimum in terms of using the available information, and for minimizing random error. At high $S/N$ and high angular resolution, however, even the most ``normal'' galaxies exhibit peculiarities (as discussed in more detail in \\S 3.3.2) such that simple bulge-plus-disk models cannot reproduce these frequently observed peculiarities in light distributions with high fidelity. The situation worsens for ``peculiar'' galaxies. For this reason we have some concerns about the uniqueness of the observationally derived model parameters, and hence their interpretation. We anticipate future developments which use the models and non-parametric measurements in a hybrid scheme optimal for characterizing galaxy light distributions both in terms of random and systematic errors.  It is worth noting again that bright galaxy samples are notorious for missing or under-representing certain galaxy types -- particularly dwarfs and low-surface-brightness galaxies. The samples used here are no exception. While this was one of our complaints about the classical Hubble scheme, there are two key differences with our approach: (i) the classification parameters we develop are objective; and (ii) these parameters do not assume the presence of basic axi-symmetry, disk-plus-bulge structure, or spiral patterns which underly the Hubble scheme. As we will show, the galaxies examined here are sufficiently diverse to establish the {\\it parameter space} for a comprehensive classification scheme, although not the comprehensive classification itself. By developing an initial classification of these galaxies, however, we intend to use it as a foil against which we can begin to compare the classification of more distant samples: How are the classifications different? Do the nearby and distant samples occupy the same regions of parameter space? If not, do the differences represent continuous extensions of these parameters, or are they physically disjoint? These are the types of questions one can address given the limitations of current local samples.  Note that we must stop short of identifying differences as ``new,'' epoch-specific classes of galaxies. Without a complete census of both the nearby and distant universe, it is not possible to establish whether there are different ``classes'' of galaxies at different redshifts; apparent differences could simply be artifacts of the presently limited samples. With such a complete census, in the future we may hope to address deeper issue of how the comoving space-densities of different classes evolve.  Towards the goal of establishing a comprehensive classification scheme of utility to distant galaxy studies, in this paper we assemble a robust set of non-parametric, photometric and structural properties for a range of nearby, lumimous galaxies. We define a multivariate, photometric parameter space that forms an initial classification scheme for these galaxies. This classification can be used reliably to identify comparable samples in other surveys and at higher redshift. In the accompanying paper (II) we measure these properties for compact, luminous emission-line galaxies at intermediate redshift, compare them to the ``normal,'' nearby galaxies studied here, and demonstrate that our classification parameter space distinguishes between these two samples. We discuss the implications for the evolution of this intermediate-redshift sample therein.  In future papers in this series we intend to extend our analysis (a) to more representative samples of the local volume that include dwarf and emission-line galaxies (e.g., the University of Michigan Objective Prism Survey (\\cite{sal89}); (b) to more comprehensive samples of distant galaxies, e.g magnitude-limited samples from the Hubble Deep Field; and (c) to studies of the morphological evolution of these distant samples. The classification scheme which we propose here is intended as a framework for these future studies.  The data sets are presented in \\S 2; the analysis is described in \\S 3. The results are presented in \\S 4, and summarized in \\S 5. Throughout this paper we adopt $H_0$ = 50 km s$^{-1}$ Mpc$^{-1}$, $q_0 = 0.1$, $\\Lambda = 0$. ",
        "conclusions": "We have presented results from a study of the photometric and image-structural characteristics and correlations of a sample of local, bright galaxies (Frei \\etal 1996). We find it illuminating to distinguish between parameters which characterize spectral-index (color), form (image concentration and asymmetry), and scale (size, luminosity, and surface-brightness). In this context, we arrive at the following main results and conclusions.  \\begin{itemize}  \\item We find that a combination of spectral-index, form and scale parameters has the greatest discriminatory power in separating normal Hubble-types. The strongest correlation is found between color and image concentration. However, there are equally strong correlations between form parameters (e.g. $A$ and $C$), but here the Hubble-types are not as well distinguished. As an indicator of classification utility, we suggest that the strength of the correlation between parameters is likely more important than the separation of Hubble Types within the correlation.  \\item It is possible to define a quantitative classification system for normal galaxies based on a four-parameter sub-set of spectral-index, form and scale: rest-frame \\bv \\ color, image concentration, asymmetry, and average surface brightness within the half-light radius.  We propose a specific classification that distinguishes between ``normal'' galaxies as Early, Intermediate, and Late based on cuts in these four parameters.  The classification is successful for 99\\% of the Frei \\etal sample.  Nonetheless, we designate this as ``preliminary'' until larger, more comprehensive samples of galaxies are needed than analyzed in the present study.  \\item Distance-dependent scale parameters are {\\it not} part of this preliminary classification.  \\item These classification parameters can be measured reliably over a broad range in $S/N$ and image resolution, and hence should be applicable to reliably distinguishing between a wide variety of galaxies over a large range in redshift.  \\item Redshift information {\\it is} needed to estimate reliably both the photometric properties (rest-frame color and surface brightness) as well as the structural parameters asymmetry and concentration at a fixed ($B$-band) rest-frame wavelength. In terms of redshift independence, asymmetry and concentration alone thus offer no advantages over the additional parameters classifiers proposed here. Indeed, incorporating the full suite of parameters defined here is advantageous for the purposes of classification.  \\end{itemize}"
    },
    "0002/astro-ph0002112_arXiv.txt": {
        "abstract": "Approximations of the Sedov self-similar solution for a strong point explosion in a medium with the power-law density distribution $\\rho^o\\propto r^{-m}$ are reviewed and their accuracy are analyzed. Taylor approximation is extended to cases $m\\neq 0$. Two approximations of the solution are presented in the Lagrangian coordinates for spherical, cylindrical and plane geometry. These approximations may be used for the investigation of the ionization structure of the adiabatic flow, i.e., inside adiabatic supernova remnants. ",
        "introduction": "\\subsection{Sedov solution}  If strong ($P_{\\rm s}/P^{o}_{\\rm s}\\rightarrow\\infty$) point ($R_{o}/R\\rightarrow0$) explosion with finite energy $E_{o}$ becomes in the point with coordinate $\\tilde{r}=0$ in time $t=0$, the blastwave creates and propagates with velocity $D$ in the ambient medium with density $\\tilde{\\rho}^{o}(\\tilde{r})$ ($P_{\\rm s}$ and $P^{o}_{\\rm s}$ are pressure of the shocked gas and the gas of the ambient medium at the shock front position, $R_{o}$ is the size of the body exploded, $R$ is the radius of the blastwave). It is also assumed that injected mass is small and no energy lost from the shocked region during the motion. \\par  Such a task is described by a system of hydrodynamic differential equations. Sedov (\\cite{Sedov}) gives the analytical self-similar solution for description of the motion of shock front and the distribution of fluid parameters inside the shocked region for a strong point explosion in the uniform ambient medium and in the center of symmetry of a radially stratified medium (\\ref{rho-power}). \\par  This solution shows that the strong blastwave in a medium with the power-law density distribution (\\ref{rho-power}) moves with deceleration then $m<N+1$ and accelerates then $m>N+1$. If $m\\geq N+1$, both the mass inside any sphere, which containes the center of the symmetry, and kinetic energy equal infinity. We will consider $m<N+1$ cases only. \\par  Radius $R$ and velocity $D$ of the strong blastwave in the medium (\\ref{rho-power}) with $m<N+1$ are (Sedov \\cite{Sedov}): \\begin{equation} \\label{R_s_Sedov-ro^w} R=\\left({E_o\\over\\alpha_A\\ \\tilde{\\rho}^o(0)}\\right)^{1/(N+3-m)} t^{\\ 2/(N+3-m)}\\ , \\end{equation} \\begin{equation} \\label{D_Sedov-ro^w} D(R)={2\\over N+3-m}\\left({E_o\\over\\alpha_A\\ \\tilde{\\rho}^o(0)} \\right)^{1/2} R^{-(N+1-m)/2}\\ , \\end{equation} \\noindent where $N=0,1,2$ for plane, cylindrical and spherical wave, respectively, $\\alpha_A$ is a self-similar constant. \\par  Distributions of the fluid characteristics behind the shock front are self-similar, i.e., for any time $t$ the density $\\tilde{\\rho}$, pressure $\\tilde{P}$, fluid velocity $\\tilde{u}$ variations and coordinate $\\tilde{a}$ are \\begin{equation} \\label{uni-prof-1} \\tilde{\\rho}(\\tilde{r},t)=\\tilde{\\rho}_{\\rm s}(t)\\cdot\\rho(r), \\end{equation} \\begin{equation} \\label{uni-prof-2} \\tilde{P}(\\tilde{r},t)=\\tilde{P}_{\\rm s}(t)\\cdot P(r), \\end{equation} \\begin{equation} \\label{uni-prof-3} \\tilde{u}(\\tilde{r},t)=\\tilde{u}_{\\rm s}(t)\\cdot u(r), \\end{equation} \\begin{equation} \\label{uni-prof-4} \\tilde{a}(\\tilde{r},t)=R(t)\\cdot a(r), \\end{equation} where $r=\\tilde{r}/R(t)$, $\\tilde{a}$ is the original position of the fluid mass element and superscript \"s\" corresponds to values of the parameters at the shock front (Fig.~\\ref{accuracy_Taylor_Kahn}). \\par  Gas occupies whole shocked region ($0\\le \\tilde{r}\\le R$) when $m\\le m_1$, \\begin{equation} m_1={1+3N+(1-N)\\gamma\\over\\gamma+1}\\ . \\end{equation} When $m\\rightarrow m_1$ central pressure $P(0)\\rightarrow 0$. Shock waves in media with steep density gradients ($m>m_1$) develop a cavity around the center of explosion. Such a cavity creates in the uniform medium ($m=0$) when $\\gamma>\\gamma_1=(1+3N)/(N-1)$. Sedov has also presented a solution for hollow blastwaves. Review of approximations for these cases is given by Ostriker \\& McKee (\\cite{Ostriker-McKee-88}). We do not consider $m>m_1$ in this paper. \\par  For $m=m_1$ (or $\\gamma=\\gamma^*$ in the uniform medium) solution has very simple form: \\begin{equation} \\begin{array}{l} \\rho(r)=r^{N-1}\\ ,\\qquad P(r)=r^{N+1}\\ , \\\\ \\\\ u(r)=r\\ ,\\hspace{1.4cm} a(r)=r^{(\\gamma+1)/(\\gamma-1)}\\ . \\end{array} \\label{solut-m_1} \\end{equation}  Singularities in the solution also appear with $m_2=(N+1)(2-\\gamma)$ and $m_3=(2\\gamma+N-1)/\\gamma$ then some exponents in the solution equal infinity. Similarity solutions for these cases are deduced by Korobejnikov \\& Rjazanov (\\cite{Korobejnikov-Rjazanov-59}). For $N=2$ and $\\gamma=5/3$ $m_1=2$, $m_2=1$, $m_3=13/5$. \\par  Self-similar constant $\\alpha_A=\\alpha_A(N,\\gamma,m)$ in equations (\\ref{R_s_Sedov-ro^w})-(\\ref{D_Sedov-ro^w}) for $R$ and $D$ may be found from the energy balance equation with variations of density $\\tilde{\\rho}$, pressure $\\tilde{P}$ and mass velocity $\\tilde{u}$ inside the shocked region \\begin{equation} \\label{enery-bal} {E_o\\over \\sigma}=\\int\\limits_0^R {\\tilde{\\rho}(r,t) \\tilde{u}(r,t)^2\\over 2} r^Ndr+ \\int\\limits_0^R {\\tilde{P}(r,t)\\over \\gamma-1} r^Ndr\\ , \\end{equation} where $\\sigma=4\\pi$ for $N=2$, $\\sigma=2\\pi$ for $N=1$ and $\\sigma=2$ for $N=0$ or, generally, $\\sigma=2\\pi N+(N-1)(N-2)$. If we proceed to normalized parameters using (\\ref{uni-prof-1})-(\\ref{uni-prof-3}) and general shock front conditions \\begin{equation} \\tilde{\\rho}_{\\rm s}={\\gamma+1\\over\\gamma-1}\\tilde{\\rho}^o_{\\rm s},\\ \\tilde{P}_{\\rm s}={2\\over\\gamma+1}\\tilde{\\rho}^o_{\\rm s}D^2,\\ \\tilde{u}_{\\rm s}={2\\over\\gamma+1}D \\end{equation} we will obtain that $E_o=\\beta_A\\cdot MD^2/2$ with \\begin{equation} M=\\sigma\\tilde{\\rho}^o(0)R^{N+1-m}/(N+1-m), \\end{equation} constant shape-factor \\begin{equation} \\beta_A={4(N+1-m)\\over\\gamma^2-1}\\cdot\\left(I_{\\rm K}+I_{\\rm T}\\right) \\label{beta_A} \\end{equation} and constant integrals \\begin{equation} I_{\\rm K}=\\int\\limits_0^1 \\rho(r)u(r)^2 r^Ndr\\ , \\qquad I_{\\rm T}=\\int\\limits_0^1 P(r) r^Ndr\\ . \\end{equation} Also we will have a self-similar constant \\begin{equation} \\label{alpha_A-vs-beta_A} \\alpha_A={2\\sigma\\over(N+1-m)(N+3-m)^2}\\cdot\\beta_A\\ . \\end{equation} Simple formula gives $\\alpha_A(N,\\gamma,m_1)$: \\begin{equation} \\alpha_A={2\\sigma(\\gamma+1)\\over(N+1)(\\gamma-1)\\big((N+1)\\gamma-N+1\\big)^2}\\ . \\end{equation}  The distributions (\\ref{uni-prof-1})-(\\ref{uni-prof-4}) in the exact solution are parametric functions of an internal parameter. The expressions for the functions are complicated. These factors stimulate developing the approximations of the self-similar solution. \\par  \\subsection{Taylor approximation}  Basing on own numerical results, Taylor (\\cite{Taylor-50}) propose to approximate the velocity variation $u(r)$ behind spherical ($N=2$) shock front moving into the uniform medium ($m=0$) as \\begin{equation} {\\tilde{u}(r,t)\\over D}={r\\over\\gamma}+\\alpha r^n, \\label{app-Taylor-base} \\end{equation} where $\\alpha$ and $n$ are found to give exact values of $\\tilde{u}_{\\rm s}$, $\\tilde{P}_{\\rm s}$, $\\tilde{\\rho}_{\\rm s}$ and their first derivatives in respect to $r$. Substituting this approximation into the continuity equation and into the equation of state for perfect gas, the approximated distributions of the density and pressure obtain. Taylor do not give the dependence $a(r)$, but it may be taken from the adiabaticity condition $P(a)\\rho(a)^{-\\gamma}=P(r)\\rho(r)^{-\\gamma}$ and (\\ref{appr_3a})-(\\ref{appr_3b}): \\begin{equation} a^{\\gamma m-(N+1)}=P(r)\\rho(r)^{-\\gamma}, \\label{a-vs-Prho} \\end{equation} with approximations for $P(r)$ and $\\rho(r)$. \\par  So, Taylor approximation for the variations of density $\\rho$, pressure $P$, fluid velocity $u$ and coordinate $a$ are: \\begin{equation} \\label{Taylor-n} \\rho(r)={\\rho(r,t)\\over \\rho_{\\rm s}(t)}= r^{\\ 3/(\\gamma-1)}\\ \\left({\\gamma+1\\over\\gamma}-{r^{n-1}\\over\\gamma}\\right)^{-p}, \\end{equation} \\begin{equation} \\label{Taylor-T} P(r)={P(r,t)\\over P_{\\rm s}(t)}= \\left({\\gamma+1\\over\\gamma}-{r^{n-1}\\over\\gamma}\\right)^{-q}, \\end{equation} \\begin{equation} \\label{Taylor-u} u(r)={u(r,t)\\over u_{\\rm s}(t)}= {\\gamma+1\\over 2}\\left({r\\over\\gamma}+{\\gamma-1\\over\\gamma+1}{r^{n}\\over\\gamma}\\right), \\end{equation} \\begin{equation} \\label{Taylor-a} a(r)={a_o(r)\\over R(t)}= r^{\\ \\gamma/(\\gamma-1)}\\ \\left({\\gamma+1\\over\\gamma}-{r^{n-1}\\over\\gamma}\\right)^{-s}, \\end{equation} where $n=(7\\gamma-1)/(\\gamma^2-1)$, $p=2(\\gamma+5)/(7-\\gamma)$, $q=(2\\gamma^2+7\\gamma-3)/(7-\\gamma)$, $s=(\\gamma+1)/(7-\\gamma)$. Self-similar constant $\\alpha_A=\\alpha_A(2,\\gamma,0)$ goes with (\\ref{alpha_A-vs-beta_A}) and approximated profiles of $\\rho$, $P$ and $u$. % Fig.~\\ref{accuracy_Taylor_Kahn} and table \\ref{alpha_comp} demonstrate accuracy of Taylor approximation in comparison with the exact solution. \\par  This approximation is extended to cases $m\\neq 0$ in section \\ref{Taylor-ext}. \\par  \\begin{figure}% \\epsfxsize=8.8truecm \\centerline{\\epsfbox{1.eps}} \\caption[]{{\\bf a-c.} Sedov solution and the accuracy of Taylor and Kahn approximations of the solution in the uniform medium: {\\bf a}~exact Sedov solution, {\\bf b}~relative differences of Taylor approximation, {\\bf c}~relative differences of Kahn approximation. Lines: 1 -- $\\rho(r)$, 2 -- $P(r)$, 3 -- $u(r)$, 4 -- $a(r)$. $\\gamma=5/3$. } \\label{accuracy_Taylor_Kahn} \\end{figure}  \\subsection{Kahn approximation}  Kahn (\\cite{Kahn}) apply his methodology to the strong spherical blastwave ($N=2$) in uniform medium ($m=0$) with $\\gamma=5/3$. It is proposed to approximate first the mass distribution \\begin{equation} \\label{mu-def} \\mu(r)={M(r,t)\\over M_{\\rm s}(t)}= 3\\int\\limits^r_0 \\rho(r)r^2dr. \\end{equation}  Sedov solution shows that $P_r(r)=0$ near the centre (subscript \"$r$\" denotes a partial derivative in respect to $r$). This fact allows to find that $\\mu_r/\\mu=15/(2r)$ at $r=0$. On the base of the equation of motion, $\\mu_r/\\mu=12$, $\\mu_{rr}=168$ and $(\\mu_r/\\mu)_r=24$ at $r=1$. Therefore ratio $\\mu_r/\\mu$ is proposed to be approximated as \\begin{equation} \\label{Kahn-dmu-appr} {\\mu_r\\over\\mu}= {15\\over 2r}+{9\\over 2}r^7. \\end{equation} This formula satisfies all written boundary conditions at both ends. \\par  The mass distribution finds as integral from (\\ref{Kahn-dmu-appr}): \\begin{equation} \\label{Kahn-mu-appr} \\mu(r)=r^{15/2}\\ \\exp\\left({9\\over16}\\left(r^8-1\\right)\\right). \\end{equation} Density distribution follows from (\\ref{mu-def}) and (\\ref{Kahn-mu-appr}): \\begin{equation} \\rho(r)=\\mu_r/3r^2. \\end{equation} Adiabaticity condition gives pressure variation \\begin{equation} P(r)=\\left({2\\over3}\\right)^{5/3}{1\\over 32}\\ {\\mu_r^{5/3}\\over \\mu\\ r^{10/3}}\\ . \\end{equation} Velocity deduses from the mass conservation equation \\begin{equation} u(r)={4\\over 3}r-{4\\mu\\over\\mu_r}\\ . \\end{equation} If present location of mass element $a$ is $r$, then a(r) may be found from the condition of mass conservation $\\mu(a)=\\mu(r)$ and relation (\\ref{mu(a)}) $\\mu(a)=a^{3}$: \\begin{equation} \\label{Kahn-a-appr} a(r)=\\mu(r)^{1/3}\\ . \\end{equation}  The expressions for Kahn approximation are the same as (\\ref{Kahn-n})-(\\ref{Kahn-mu}) with $m=0$. The accuracy of this approximation are shown on the Fig.~\\ref{accuracy_Taylor_Kahn}.  \\subsection{Approximation of Cox \\& Franco}  Appling Kahn's approximation technique, Cox \\& Franco (\\cite{Cox_Fanko-81}) obtain the approximation of the self-similar solution for an ambient medium with the power-law density distribution (\\ref{rho-power}) with $m<2$ for $\\gamma=5/3$ and $N=2$. Approximation of Cox \\& Franco are: \\begin{equation} \\begin{array}{l} {\\displaystyle \\rho(r)=% \\left({5\\over8}+{3\\over8}r^{\\ 8-4m}\\right)\\cdot r^{\\ (9-5m)/2} }\\\\ \\\\ {\\displaystyle \\qquad\\qquad\\qquad\\!\\! \\times\\exp\\left({3\\over8}{(3-m)\\over(2-m)}\\left(r^{\\ 8-4m}-1\\right)\\right)\\ , } \\end{array} \\label{Kahn-n} \\end{equation} \\begin{equation} \\begin{array}{l} {\\displaystyle P(r)=% \\left({5\\over8}+{3\\over8}r^{\\ 8-4m}\\right)^{5/3} }\\\\ \\\\ {\\displaystyle \\qquad\\qquad\\qquad\\!\\! \\times\\exp\\left({3\\over4(2-m)}\\left(r^{\\ 8-4m}-1\\right)\\right)\\ , } \\end{array} \\label{Kahn-T} \\end{equation} \\begin{equation} \\label{Kahn-u} u(r)=% 4\\ r\\ {1+r^{\\ 8-4m}\\over 5+3\\ r^{\\ 8-4m}}\\ , \\end{equation} \\begin{equation} {\\displaystyle a(r)=% r^{5/2}\\ \\exp\\left({3\\over8(2-m)}\\left(r^{\\ 8-4m}-1\\right)\\right)\\ ,} \\label{Kahn-a} \\end{equation} \\begin{equation} {\\displaystyle \\mu(r)=r^{\\ 5(3-m)/2} \\ \\exp\\left({3\\over8}{(3-m)\\over(2-m)}\\left(r^{\\ 8-4m}-1\\right)\\right)\\ . } \\label{Kahn-mu} \\end{equation} Author's approximation for $\\beta_A$ is \\begin{equation} \\label{beta-Cox-Fanko} \\beta_A=1.125\\cdot(0.22+0.52\\cdot(3-m)/3). \\end{equation}  The accuracy of Cox \\& Franco approximation is shown on Fig.~\\ref{accuracy_Kahn_N-2} and table \\ref{alpha_comp}. \\par  \\begin{figure}% \\epsfxsize=8.8truecm \\centerline{\\epsfbox{2.eps}} \\caption[]{{\\bf a-c.} Accuracy of Cox \\& Franco approximation of the self-similar solution in the power-law medium (\\ref{rho-power}): {\\bf a}~relative differences of the approximation for $m=-4$, {\\bf b}~relative differences for $m=-2$, {\\bf c}~relative differences for $m=1$. Lines are the same as on Fig.~\\ref{accuracy_Taylor_Kahn}. Fracture in the curves for $\\rho(r)$ and $a(r)$ is due to very strong dependence of the relevant Sedov distributions on the internal parameter, which changes in these wide intervals of $r$ on $10^{-10}$ only. } \\label{accuracy_Kahn_N-2} \\end{figure}  \\subsection{Approximations of Ostriker \\& McKee} \\label{Ostr-McKee}  Ostriker \\& McKee (\\cite{Ostriker-McKee-88}) in the frame of the virial theorem approach applied to spherical blastwave (N=2) in the power-law ambient medium (\\ref{rho-power}) and time-dependent energy injection $E_o(t)\\propto t^{s}$, present a number of approximations for the self-similar solution. We consider further $s=0$. \\par  Authors introduce the dimensionless moments of coordinate $r$ and velocity $u$: \\begin{equation} \\label{moment-Ostr-McKee} K_{ij}=l_\\mu\\int\\limits_0^1 r^iu(r)^j\\rho(r)r^2dr\\ , \\end{equation} where $l_\\mu=(\\gamma+1)(3-m)/(\\gamma-1)$, and consider three types of approximations for $u(r)$ and $\\rho(r)$: linear velocity approximation (LVA) \\begin{equation} \\label{LVA-Ostr-McKee} u(r)=r,\\qquad \\rho(r)=r^{(6-(\\gamma+1)m)/(\\gamma-1)}, \\end{equation} one-power aproximation (OPA) \\begin{equation} \\label{OPA-Ostr-McKee} u(r)=r^{l_u},\\qquad \\rho(r)=r^{l_\\rho}, \\end{equation} and two-power aproximation (TPA) \\begin{equation} \\label{TPA-Ostr-McKee-u} u(r)=a_ur^{l_{u,1}}+(1-a_u)r^{l_{u,2}},\\\\ \\\\ \\end{equation} \\begin{equation} \\label{TPA-Ostr-McKee-rho} \\rho(r)=a_\\rho r^{l_{\\rho,1}}+(1-a_\\rho)r^{l_{\\rho,2}}. \\end{equation}  In such an approach the self-similar constant $\\alpha_A$ as well as exponents $l_u$ and $l_\\rho$ may be expressed in terms of moments $K_{02}$ and $K_{11}$. Namely, under self-similarity $\\alpha_A=2\\pi\\eta^2\\beta_A/(3-m)$, where $\\eta=2/(5-m)$ and factor $\\beta_A$ equals \\begin{equation} \\label{beta_A-Ostr-McKee} \\beta_A={2\\over3}\\cdot{2K_{02}(3\\gamma-5)+(5-m)(\\gamma+1)K_{11} \\over (\\gamma^2-1)(\\gamma+1)}\\ . \\end{equation} Exponents in OPA are \\begin{equation} l_u={2K_{20}-K_{11}(1+K_{20})\\over(1-K_{20})K_{11}}\\ , \\quad l_{\\rho}={5K_{20}-3\\over1-K_{20}}\\ . \\end{equation}  Derivatives at shock front are used to obtain the moments. So, \\begin{equation} K_{ij}={1\\over 1+s_{ij}/l_\\mu}\\ , \\label{K_Ostr-McKee} \\end{equation} where $s_{ij}=i+j$ in LVA and $s_{ij}=i+j+j(m_1-m)/2$ in OPA. Using (\\ref{K_Ostr-McKee}) $\\alpha_A$ may be written in a simple form in LVA: \\begin{equation} \\label{alpha_A-Ostr-McKee} \\alpha_A={16\\pi\\over 3(5-m)^2}\\cdot {11\\gamma-5-m(\\gamma+1)\\over (\\gamma^2-1)\\big(5\\gamma+1-m(\\gamma+1)\\big)}\\ . \\end{equation}  Moments have more complicated form in TPA. In this approach the expression for $u(r)$ coinsides with the approximation (\\ref{Taylor-u}) of Taylor with $n=1+\\gamma(m_1-m)/(\\gamma-1)$ that equals to Taylor's $n$ at $m=0$. So, TPA is extension of Taylor approximation of $u(r)$ to $m\\neq 0$. Contrary to Taylor's approach to find $\\rho(r)$ and $P(r)$ from the hydrodynamic equations, Ostriker \\& McKee find the density variation independently as TPA (\\ref{TPA-Ostr-McKee-rho}) with \\begin{equation} \\begin{array}{l} {\\displaystyle a_{\\rho}={\\gamma(m_1-m)\\over 10-\\gamma-(\\gamma+2)m} \\ ,}\\\\ \\\\ {\\displaystyle l_{\\rho,1}={3-\\gamma m\\over\\gamma-1} \\ , \\quad l_{\\rho,2}={6+(\\gamma+1)(m_1-2m)\\over \\gamma-1} \\ ,} \\end{array} \\end{equation} where $\\gamma>1$. For $m=0$ variation $\\rho(r)$ in TPA coinsides with the result of Gaffet (\\cite{Gaffet}) for case of uniform medium (Ostriker \\& McKee \\cite{Ostriker-McKee-88}). Two-power velocity approximation is used to extend Taylor approximation to cases $m\\neq 0$ in section \\ref{Taylor-ext}. \\par  Pressure distribution are also restored independently. It may be found in OPA as a linear pressure approximation (LPA) and for TPA in the frame of pressure-gradient approximation (PGA). \\par  Most of mass is concentrated near the shock front and distribution $u(r)$ is close to a linear function of $r$. Therefore, as noted by Gaffet (\\cite{Gaffet}), the right side of Euler equation \\begin{equation} {\\partial\\tilde{P}(\\tilde{r},t)\\over \\partial M(\\tilde{r},t)}= -{1\\over4\\pi}{1\\over \\tilde{r}^2}{d\\tilde{u}(\\tilde{r},t)\\over dt} \\label{Eul-eq} \\end{equation} is nearly a constant. LPA (Gaffet \\cite{Gaffet}, Ostriker \\& McKee \\cite{Ostriker-McKee-88}) use this feature assuming the pressure to be a linear function of the mass fraction $\\mu(r)$ \\begin{equation} P(r)=P(0)+(P^*_{\\rm s}/l_{\\mu})\\ \\mu(r). \\end{equation} Logariphmic derivative of pressure at the shock front is $P^*_{\\rm s}=(d\\ln P/d\\ln r)_{\\rm s}= (2\\gamma^2+7\\gamma-3-\\gamma m(\\gamma+1))/(\\gamma^2-1)$. Mass in OPA is $\\mu(r)=3l_\\mu^{-1}r^{l_\\mu}$. P(0) in LPA is (Gaffet \\cite{Gaffet}) \\begin{equation} P(0)=1+{\\overline{u_t^{\\rm s}}\\over\\omega(3-m)} \\end{equation} where $\\overline{u_t^{\\rm s}}=\\tilde{u}_t^{\\rm s}R/D^2= \\omega\\big((4-3\\omega)(m-3)+2(1-\\omega)(4-2\\omega-m)\\big)/2$ (Hnatyk \\cite{Hn87}), $\\omega=2/(\\gamma+1)$. \\par  Such an approach (substitution with $\\tilde{r}^{-2}_{\\rm s}\\tilde{u}_t^{\\rm s}$ instead of $\\tilde{r}^{-2}\\tilde{u}_t$ in (\\ref{Eul-eq})) was also used by Laumbach \\& Probstein (\\cite{L-P}) to develop the sector approximation. \\par  In PGA a power-law form for the pressure gradient \\begin{equation} {dP(r)\\over dr}=P^*_{\\rm s}r^{l_{p,2}-1}\\ \\end{equation} is used to give two-power approximation for the pressure \\begin{equation} P(r)=P(0)+a_pr^{l_{p,2}}\\ , \\end{equation} where $a_p=P^*_{\\rm s}/l_{p,2}$ and \\begin{equation} \\begin{array}{l} {\\displaystyle P(0)={(\\gamma+1)^2(m_1-m)\\over 3\\gamma^2+20\\gamma+1-(\\gamma+1)(3\\gamma+1)m}\\ ,}\\\\ \\\\ {\\displaystyle \\hspace{0.4cm} l_{p,2}= {3\\gamma^2+20\\gamma+1-(\\gamma+1)(3\\gamma+1)m \\over 2(\\gamma^2-1)} \\ .} \\end{array} \\end{equation}  Accuracy in determination of $\\alpha_A$ and $P(0)$ in approximations of Ostriker \\& McKee is shown in table \\ref{alpha_comp} and in revealing the flow parameters on Fig.~\\ref{accuracy_Ostr-McKee}. \\par  \\begin{figure}% \\epsfxsize=8.6truecm \\centerline{\\epsfbox{3.eps}} \\caption[]{Relative differences of Ostriker \\& McKee two-power approximation of the self-similar solution for the uniform medium. $\\gamma=5/3$. Lines are the same as on Fig.~\\ref{accuracy_Taylor_Kahn}. } \\label{accuracy_Ostr-McKee} \\end{figure}  \\subsection{Cavaliere \\& Messina approximation of $\\alpha_A$}  Cavaliere \\& Messina (\\cite{Cavaliere-Messina-76}) with a simple technique approximate the equations for the radius and velocity of shock in the power-law medium (\\ref{rho-power}) and $E_o(t)\\propto t^{s}$. For $s=0$ his approximation gives \\begin{equation} \\label{beta-Cav-Mess} \\beta_A={4\\over \\gamma^2-1}\\left({\\gamma-1\\over\\gamma+1}+{1\\over 2} {N+1-m\\over N+1}\\right). \\end{equation}  \\subsection{Approximate methods for an explosion in medium with arbitrary large-scale nonuniformity}  In this subsection we pointed out a number of approximate methods for description of a point explosion in arbitrary nonuniform medium. These methods may also be applicable for a medium with power-law density variation. Bisnovatyi-Kogan \\& Silich (\\cite{BK-Syl}) and Hnatyk (\\cite{Hn87}) have given the reviews of these methods, their applications and accuracy.  \\subsubsection{Thin-layer approximation}  Thin-layer approximation is firstly introduced by Chernyi (\\cite{Chernyi}) and used by Kompaneets (\\cite{Komp}) and other authors to find analytical solutions for evolution of the shock front in a number of type of nonuniform media. It is assumed in this approach that all swept-up mass is concentrated in the infinitely thin layer just after shock front and the motion is stimulated with the hot gas inside the shocked region with uniform pressure distribution $P(r)=0.5$ (excepting $P_{\\rm s}=1$). Layer of the gas moves with velocity $u_{\\rm s}$. This method was developed to calculate anly the shock front dynamics and therefore does not allow to reveal the distribution of the fluid parameters behind the shock front. \\par  Thin-layer approximation gives for spherical blastwave in the uniform medium (Andriankin et al. \\cite{Andriankin-et-al}) \\begin{equation} \\label{alpha_A-TL} \\alpha_A={16\\pi(3\\gamma-1)\\over 75(\\gamma-1)(\\gamma+1)^2}\\ . \\end{equation}  \\subsubsection{Sector approximation}  In the sector approximation, the characteristics of an one-dimentional flow find as decompositions into series about the shock front. \\par  Laumbach \\& Probstein (\\cite{L-P}) have proposed the sector approximation applying it to spherical blastwaves in a plane-stratified exponential medium. Authors use Lagrangian coordinate $a$ and propose to approximate pressure variation in the form equivalent to $P(a)=1+P_a^{\\rm s}(a-1)$ (Hnatyk \\cite{Hn87}). Density variation is given by the adiabaticity condition and relation $r=r(a)$ by continuity equation. Fluid velocity field is not determined. For shock radius and its velocity Laumbach \\& Probstein approximation yeilds in the uniform medium limit \\begin{equation} \\label{alpha_A-SA} \\alpha_A={32\\pi(4\\gamma^2-\\gamma+3)\\over 225(\\gamma-1)(\\gamma+1)^3}\\ . \\end{equation}  Gaffet (\\cite{Gaffet,Gaffet-81}) uses Lagrangian mass coordinates $\\mu$ and finds pressure variation as a linear pressure approximation $P(\\mu)=1+P_\\mu^{\\rm s}(\\mu-1)$. Gaffet (\\cite{Gaffet,Gaffet-81}) also propose to improve accuracy of the approximation, taking into account the second order coefficients in the series. Author calculates such coefficients in terms of Lagrangian mass coordinate $\\mu$. Hnatyk (\\cite{Hn87}), considering different modifications of the sector approximation, presents the coefficients up to the second order in terms of $a$. \\par  \\begin{figure}% \\epsfxsize=8.6truecm \\centerline{\\epsfbox{4.eps}} \\caption[]{Accuracy of Hnatyk approximation of Sedov solution for the uniform medium. Lines: 1 -- $\\rho(a)$, 2 -- $P(a)$, 3 -- $u(a)$, 4 -- $r(a)$. $\\gamma=5/3$. } \\label{accuracy_Hn} \\end{figure}  \\subsubsection{Hnatyk approximation} \\label{Hn-app}  Hnatyk (\\cite{Hn87}) introduces also the idea to aproximate first\\-ly the relation $\\tilde{r}=\\tilde{r}(a,t)$ between the Lagrangian $a$ and Eulerian $r$ coordinates of the gas element in each sector of shocked region. Density $\\rho$, pressure $P$ and velocity $u$ variation behind the shock front are exactly deduced from this relation. Really, the continuity equation \\begin{equation} \\tilde{\\rho}^o(\\tilde{a}) \\tilde{a}^N d\\tilde{a}= \\tilde{\\rho}(\\tilde{r}) \\tilde{r}^N d\\tilde{r} \\label{cont-eq-rho} \\end{equation} gives us the density distribution \\begin{equation} \\label{ro(at)} \\rho(a)={\\tilde{\\rho}(a,t)\\over \\tilde{\\rho}^{\\rm s}(t)}= {\\tilde{\\rho}^{o}(\\tilde{a})\\over \\tilde{\\rho}(R,t)} \\left({\\tilde{a}\\over \\tilde{r}(\\tilde{a},t)}\\right)^N \\left({\\partial \\tilde{r}(\\tilde{a},t)\\over\\partial \\tilde{a}} \\right)^{-1}, \\end{equation} the equation of adiabaticity \\begin{equation} \\tilde{P}(\\tilde{a},t)=K\\tilde{\\rho}(\\tilde{a},t)^{\\gamma} \\end{equation} yields the distribution of pressure \\begin{equation} \\label{P(at)} P(a)= {\\tilde{P}(\\tilde{a},t)\\over \\tilde{P}^{\\rm s}(t)}= \\biggl({\\tilde{\\rho}^{o}(\\tilde{a})\\over \\tilde{\\rho}^{o}(R)}\\biggr)^{\\!1-\\gamma} \\biggl({D(\\tilde{a})\\over D(R)}\\biggr)^{\\!2}\\biggl({\\tilde{\\rho}(\\tilde{a},t)\\over \\tilde{\\rho}(R,t)} \\biggr)^{\\!\\gamma} \\end{equation} and relation $\\tilde{r}=\\tilde{r}(\\tilde{a},t)$ gives velocity \\begin{equation} \\label{u(at)} u(a)= {\\tilde{u}(\\tilde{a},t)\\over \\tilde{u}^{\\rm s}(t)}= {\\gamma+1\\over 2}{1\\over D(R)} {d\\tilde{r}(\\tilde{a},t)\\over dt}\\ . \\end{equation}  Author propose to approximate $r(a)$ as \\begin{equation} r(a)=a^\\alpha\\exp\\big(\\beta(a-1)\\big)\\ \\label{r(a)-Hn} \\end{equation} with \\begin{equation} \\alpha=(r^{\\rm s}_{a})^2-r^{\\rm s}_{aa} \\quad {\\rm and}\\quad \\beta=r^{\\rm s}_{aa}+r^{\\rm s}_{a}-(r^{\\rm s}_{a})^2 \\ . \\end{equation} Such an expression ensures the edge condition $r(0)=0$, $r_{\\rm s}=1$ and values of the derivatives \\begin{equation} r^{\\rm s}_a = 1-\\omega, \\label{ras} \\end{equation} \\begin{equation} r^{\\rm s}_{aa}=\\omega(1-\\omega) \\bigl[ 3B+N(2-\\omega)-m\\bigr] \\label{raas} \\end{equation} where $B=R\\ddot R/\\dot R^2$, $\\dot R=dR/dt$ is the shock velocity, $m=-d\\ln\\rho^o(R)/d\\ln R$, subscript \"$a$\" denotes a partial derivative in respect to $a$. \\par  This approximation is accurate near the shock front, but around the explosion site (for $a<0.1$ or $r<0.4$) characteristics do not restore correctly (Fig.~\\ref{accuracy_Hn}). This approximation does not take into consideration any derivatives of $r(a)$ near the center and the distributions of $\\rho(a)$, $P(a)$, $u(a)$ do not bind there, causing such a situation. This approximation is extended to the central region in subsection \\ref{app-Hn-imp}. \\par ",
        "conclusions": "In this paper, we review approximations of the self-similar solution for a strong point explosion in the power law medium $\\rho^o\\propto r^{-m}$ and compare their accuracy with the exact Sedov solution of the problem. Different approaches found on the different basic approximations. Namely, Taylor (\\cite{Taylor-50}) and Ostriker \\& McKee (\\cite{Ostriker-McKee-88}) approximate first\\-ly the fluid velocity variation behind the shock front. Taylor used approximated $u(r)$ substituting it into the hydrodynamic equations to obtain full description of the flow. Contrary to this, Ostriker \\& McKee approximate $\\rho(r)$ and $P(r)$ independently. Kahn (\\cite{Kahn}) technic, used also by Cox \\& Franco (\\cite{Cox_Fanko-81}), consists in approximation of the fluid mass variation $\\mu(r)$ and further usage of the system of hydrodynamic equations. Gaffet (\\cite{Gaffet}), Laumbach \\& Probstein (\\cite{L-P}), Ostriker \\& McKee (\\cite{Ostriker-McKee-88}) base their approaches on the approximation of $P(\\mu)$ or $P(r)$. Thin layer approximation may also be included into this group. Hnatyk (\\cite{Hn87}) take approximation of the connection between Eulerian and Lagran\\-gi\\-an coordinates as basic relation. So, practically all possible approaches are used to have approximation for the self-similar solution. \\par  In this paper we apply Taylor's methodology to discribe a strong point explosion in the power-law medium, extending his approximation written for uniform medium, and write also two approximations expressed in Lagran\\-gi\\-an geometric coordinates, approaching $r(a)$ with different accuracy. \\par  Errors of all approximations are caused only by errors in the basic approximation. When the first approximation has higher accuracy we have more accurate approximation for parameters of the shock and flow. \\par  \\begin{table*}% \\caption[]{ Comparision of the self-similar constant $\\alpha_{\\rm A}$ and pressure $P(0)$ calculated with: S --  Sedov (\\cite{Sedov-1946b}) solution (Kestenboim et al. \\cite{Kestenb}); T --  Taylor (\\cite{Taylor-50}) approximation; CF -- approximation of Cox \\& Franco (\\cite{Cox_Fanko-81}); LVA, OPA and TPA of Ostriker \\& McKee (\\cite{Ostriker-McKee-88}); CM -- approximation of Cavaliere \\& Messina (\\cite{Cavaliere-Messina-76}); TL -- thin-layer (\\ref{alpha_A-TL}) approximation; LP -- approximation (\\ref{alpha_A-SA}) of Laumbach \\& Probstein (\\cite{L-P}); SOA -- second order (\\ref{Hn-impruved}) and TOA -- third order (\\ref{r(a)_approx}) approximations. Uniform medium, $\\gamma=5/3$ and $N=2$. \\\\ } \\begin{center} \\begin{tabular}{lcccccccccccccc} \\hline \\noalign{\\smallskip} &S&T&CF&LVA&OPA/LPA&TPA/PGA&CM&TL&LP&SOA&TOA\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} $\\alpha_A$&0.4936 &0.4957 &0.4930 &0.5386 &0.5027 &0.4957  &0.5655 &0.5655 &0.4398 &0.4981  &0.4944 \\\\ \\noalign{\\smallskip} $P(0)$\t  &0.3062 &0.2855 &0.3140 &--     &0.3333 &0.3333  &--     &0.5000 &0.3333 &0.2507  &0.3062 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\label{alpha_comp} \\end{table*}   \\appendix"
    },
    "0002/astro-ph0002438_arXiv.txt": {
        "abstract": "We have determined the radial velocities of the [O~III] emitting gas in the inner narrow-line region (NLR) of the Seyfert 2 galaxy NGC 1068, along a slit at position angle 202\\deg, from STIS observations at a spatial resolution of 0\\arcsecpoint1 and a spectral resolving power of $\\lambda$/$\\Delta\\lambda$ $\\approx$ 1000. We use these data to investigate the kinematics of the NLR within 6$''$ ($\\sim$430 pc) of the nucleus. The emission-line knots show evidence for radial acceleration, to a projected angular distance of 1\\arcsecpoint7 in most cases, followed by deceleration that approaches the systemic velocity at a projected distance of $\\sim$4$''$. We find that a simple kinematic model of biconical radial outflow can match the general trend of observed radial velocities. In this model, the emitting material is evacuated along the bicone axis, and the axis is inclined 5\\deg\\ out of the plane of the sky. The acceleration of the emission-line clouds provides support for dynamical models that invoke radiation and/or wind pressure. We suggest that the deceleration of the clouds is due to their collision with a patchy and anistropically distributed ambient medium. ",
        "introduction": "The narrow-line region (NLR) in Seyfert galaxies is characterized by emission lines with widths on the order of 500 km s$^{-1}$ (full-width at half-maximum, FWHM), which are attributed to motions of the ionized clouds of gas. Over the past couple of decades, ground-based studies attempted to determine the kinematics of the NLR, but a general consensus on the velocity flow pattern was not reached; cases were made for infall, rotation, parabolic orbits, outflow, etc. (e.g., Osterbrock and Mathews 1986, and references therein; DeRobertis \\& Shaw 1990; Veilleux 1991; Moore \\& Cohen 1994, 1996). Since the majority of the NLR flux comes from a region that subtends only a few arcseconds in most Seyferts (Schmitt \\& Kinney 1996), these studies had to rely on spatially-integrated line profiles. Unfortunately, similar profile shapes and asymmetries can be generated from a wide variety of kinematic models (Capriotti, Foltz, \\& Byard 1980, 1981), and hence the difficulty in determining the velocity fields from these data. By contrast, ground-based studies of the extended narrow-line region (ENLR, at distances typically $\\geq$ 500 pc from the central source) can take advantage of of spatially-resolved spectra; these studies find that the ionized gas in the ENLR is undergoing normal galactic rotation (Unger et al. 1987), with evidence for an additional component of outward radial motion in some cases (Whittle et al. 1988). Despite the limited spatial resolution, recent ground-based studies suggest that the NLR of NGC 1068, the nearest bright Seyfert 2 galaxy, shows evidence for radial outflow (Cecil, Bland, \\& Tully 1990; Arribas, Mediavilla, and Garc\\'{i}a-Lorenzo 1996), which is a suggestion first offered by Walker (1968).  With the {\\it Hubble Space Telescope} (HST) and the Space Telescope Imaging Spectrograph (STIS), we now have the ability to obtain spectra of the NLR at high spatial resolution. The importance of these observations is that we can probe the velocity field of the ionized gas close to the central continuum source, where the supermassive black hole presumably dominates the kinematics (due to its gravitational influence and/or the radiation, winds, and jets emanating from its vicinity). In this letter, we use STIS long-slit spectra of the Seyfert 2 galaxy NGC 1068 to determine the kinematics of the ionized gas in its NLR. In previous papers, we used these spectra to study the extended continuum emission (Crenshaw \\& Kraemer 2000, hereafter Paper I) and the physical conditions in ionized gas near the continuum ``hot spot'' (Kraemer \\& Crenshaw 2000, Paper II). We adopt a systemic redshift of cz $=$ 1148 km~s$^{-1}$ from H~I observations of NGC 1068 (Brinks et al. 1997) and a distance of 14.4 Mpc (Bland-Hawthorne 1997), so that 0\\arcsecpoint1 corresponds to 7.2 pc. ",
        "conclusions": "\\subsection{Comparison with Ground-based Observations}  Cecil et al. (1990) provide the most comprehensive set of velocity maps for the NLR of NGC 1068, based on Fabry-Perot observations of the [N II] lines at $\\sim$1$''$ spatial and 140 km s$^{-1}$ spectral resolutions; these authors conclude that their data can be explained by cylindrical or biconical outflow. Given the limited spatial resolution, their observations are in agreement with ours and are consistent with our model. Their line profiles along the STIS slit position show a single broad component at the location of the hot spot and double-peaked profiles, separated by $\\sim$1000 km s$^{-1}$, at a distance $\\sim$2$''$ from the hot spot, in agreement with our observations (Figure 3). Outside of the STIS slit position, their profiles show the same double-peaked structure, with the largest velocity separation at $\\sim$2$''$ on either side of the hot spot, which is consistent with our model prediction that the highest blueshifts and redshifts should occur at this distance. The higher spatial resolution of the STIS observations allows us to see the acceleration of clouds outward from the nucleus followed by a clear deceleration of the clouds. We have not applied our model to ground-based observations at distances $>$ 6$''$, where galactic rotation begins to play a much larger role in the kinematics (Baldwin, Whittle, \\& Wilson 1987).  \\subsection{Comparison with Other Models}  Two other kinematic models of the NLR have been proposed on the basis of spectra at high spatial resolution, which were obtained with {\\it HST}'s Faint Object Camera (FOC). Axon et al. (1998) suggest a model for NGC 1068 in which the gas expands outward from the radio jet (which is nearly coincident with the axis of the ionization cone). Winge et al. (1999) propose a rotating disk model for the NLR in the Seyfert 1 galaxy NGC~4151. To test the Axon et al. (1998) suggestion for NGC 1068, we generated a model with the same parameters as in Table 1, except that the velocity vectors are directed perpendicular to the radio axis. In this case, we find that the envelope of radial velocities is nearly the same as in Figure 3, due to the small inclination angle. However, we have two concerns about this model. First, the observed radial velocities follow a well-organized flow pattern, and we can discern no correlation with the clumpy radio structure in the NLR (Gallimore et al. 1996). Second, this model cannot explain the kinematics of the NLR in the Seyfert 1 galaxy NGC 4151, where the gas is blueshifted in one cone and redshifted in the other (Winge et al. 1999; Nelson et. al. 2000, and references therein), whereas motions perpendicular to the axis produce equal blueshifts and redshifts in each cone, regardless of the inclination of the axis. We note that our radial outflow model for NGC 1068 provides a slightly better fit than the Axon et al. model, because a small tilt of the axis (5\\deg) can match the slightly different amplitudes of the observed blueshifted and redshifted curves. Furthermore, by tilting the cone axis towards the line of sight, this model can explain the observed radial velocities in NGC 4151 (see Crenshaw et al. 2000).  Winge et al.'s (1999) rotating disk model for NGC 4151 is only used to match their low velocity component (within 300 km s$^{-1}$ of systemic). Even so, they require an extended (and otherwise undetected) distribution of matter within 0\\arcsecpoint1 (64 pc) of the nucleus with a mass on the order of 10$^{9}$ M$\\odot$, in addition to a central point-source mass of 10$^{7}$M$\\odot$. A rotation model can be ruled out for NGC 1068, because high redshifts and blueshifts are seen on either side of the nucleus. Other gravitational models can also be ruled out as the principal source of the velocities in NGC 1068, because the required mass is prohibitive. At the peaks of the velocity curves (ignoring the two knots with very high velocities), the projected distance from the nucleus is $\\sim$120 pc, the projected velocity is $\\sim$850 km s$^{-1}$, and the required mass is $\\geq$ 10$^{10}$ M$\\odot$ (a lower limit because of projection effects). By comparison, the dynamical mass from stars within a radius of $\\sim$1$''$ from the nucleus of NGC 1068 is only 6 x 10$^{8}$ M$\\odot$ (Thatte et al. 1997). Thus, radial outflow provides the simplest and best explanation of the observed velocities in these two Seyfert galaxies.  \\subsection{Implications of our Model}  Our kinematic model assumes constant acceleration of clouds in the inner NLR ($<$ 140 pc), and constant deceleration further out; this assumption provides a reasonable match to the observations, although more complicated velocity laws as a function of distance are possible, given the intrinsic scatter in the observed points. Nevertheless, these results favor dynamical models that invoke radiation and/or wind pressure to drive clouds out from the nucleus. The deceleration of clouds further is not primarily due to gravity, for the reasons given above: an unreasonably high mass ($\\sim$10$^{10}$ M$\\odot$) is required to slow the clouds down. The simplest explanation for the deceleration is that the clouds experience a drag force, presumably due to interaction with ambient material at $\\sim$140 pc from the nucleus. A possible explanation for the blueshifted clouds on the SW side is that they run into ambient material that is closer to the nucleus ($\\sim$80 pc). In this picture, the two high velocity clumps in Figure 3 represent clouds that have not experienced a drag force in the direction they are traveling, which suggests that there are holes in the surrounding medium. We note that gravity may eventually play a role in the kinematics. In our model, the axis of the outflow is inclined by $\\sim$45\\deg\\ with respect to the galactic disk (inclination $=$ 40\\deg, major axis at position angle 106 \\deg, see Bland-Hawthorne et al. 1997), and as the NLR clouds slow down, they may be pulled back to the disk, possibly joining the existing ENLR gas.  In conclusion, we find that a biconical outflow model, with evidence for acceleration close to the nucleus and deceleration further out, provides a reasonable explanation for the radial velocities in our long-slit spectrum of NGC 1068. STIS long-slit observations of NGC 1068 at higher spectral resolution and at different slit positions will be helpful in resolving velocity components and mapping out the two-dimensional velocity field. STIS observations of other Seyferts will help test the utility of this model."
    },
    "0002/astro-ph0002324_arXiv.txt": {
        "abstract": "We present a preliminary analysis of the 1--10 keV spectrum of the massive X-ray binary Cyg X-3, obtained with the High Energy Transmission Grating Spectrometer on the {\\it Chandra X-ray Observatory}. The source reveals a richly detailed discrete emission spectrum, with clear signatures of photoionization-driven excitation. Among the spectroscopic novelties in the data are the first astrophysical detections of a number of He-like 'triplets' (Si, S, Ar) with emission line ratios characteristic of photoionization equilibrium, fully resolved narrow radiative recombination continua of Mg, Si, and S, the presence of the H-like Fe Balmer series, and a clear detection of a $\\sim 800$ km s$^{-1}$ large scale velocity field, as well as a $\\sim 1500$ km s$^{-1}$ FWHM Doppler broadening in the source. We briefly touch on the implications of these findings for the structure of the Wolf-Rayet wind. ",
        "introduction": "In a previous paper (Liedahl \\& Paerels 1996, 'LP96') we presented an interpretation of the discrete spectrum of Cyg X-3 as observed with the Solid State Imaging Spectrometers on {\\it ASCA} (cf. Kitamoto et al. 1994; Kawashima \\& Kitamoto 1996). We found clear spectroscopic evidence that the discrete emission is excited by recombination in a tenuous X-ray photoionized medium, presumably the stellar wind from the Wolf-Rayet companion star (van Kerkwijk et al. 1992). Specifically, the {\\it ASCA} spectrum revealed a narrow radiative recombination continuum (RRC) from H-like S, unblended with any other transitions. On closer inspection, RRC features due to H-like Mg  and Si were also found to be present in the data, although severely blended with emission lines. These narrow continua are an unambiguous indicator of excitation by recombination in X-ray photoionized gas, and their relative narrowness is a direct consequence of the fact that a highly ionized photoionized plasma is generally much cooler than a collisionally ionized plasma of comparable mean ionization (LP96, Liedahl 1999 and references therein).  With the high spectral resolution of the {\\it Chandra} High Energy Transmission Grating Spectrometer, we now have the capability to fully resolve the discrete spectrum. Apart from offering a unique way to determine the structure of the wind of a massive star, study of the spectrum may yield other significant benefits. Cyg X-3 shows a bright, purely photoionization driven spectrum, and, as such, may provide a template for the study of the spectra of more complex accretion-driven sources, such as AGN. The analysis will also allow us to verify explicitly the predictions for the structure of X-ray photoionized nebulae derived from widely applied X-ray photoionization codes. ",
        "conclusions": "The HETGS spectrum of Cyg X-3 has revealed a rich discrete spectrum, the properties of which are consistent with pure recombination excitation in cool, optically thin, low density X-ray photoionized gas in equilibrium. We fully resolve the narrow RRCs for the first time, and estimate an average electron temperature in the photoionized region of $kT_e \\sim 50$ eV, consistent with global photoionization calculations.  We detect a net redshift in the emission lines of $v \\sim 750-800$ km s$^{-1}$, essentially independent of binary phase, and a distribution in velocity with a FWHM of $\\sim 1500$ km s$^{-1}$. If the wind were photoionized throughout, we would expect to see roughly equal amounts of blue-- and redshifted material, so evidently we are viewing an ionized region that is not symmetric with respect to the source of the wind, as expected if only the part of the wind in the vicinity of the X-ray continuum source is ionized. However, in the simplest wind models, one would then expect to see a strong dependence of the centroid velocity on binary phase, alternating between red-- and blueshifts, and this is decidedly not the case in our data. The implications of this finding for the flow pattern and distribution of material in the wind will be explored in a future paper.  Finally, the Fe K$\\alpha$ fluorescent feature, which probes a more neutral phase of the wind, has never been seen before in Cyg X-3. Unfortunately, the exact range of ionization can not be separated uniquely from systematic Doppler shifts through a measurement of the wavelengths of the K$\\alpha$ spectra, because the feature, while clearly broadened, is not separated into its component ionization stages. Still, the width of the feature (the net effect of the velocity field and the existence of a range of charge states) and its intensity will impose strong constraints on the global properties of the wind.   \\noindent Acknowledgements. \\newline \\noindent We wish to express our gratitude to Dan Dewey and Marten van Kerkwijk, for discussions and a careful reading of the manuscript, and to the referee, Randall Smith, for a thorough review. JC acknowledges support from NASA under a GRSP fellowship. MS's contribution was supported by NASA under Long Term Space Astrophysics grant no. NAG 5-3541. FP was supported under NASA Contract no. NAS 5-31429. DL acknowledges support from NASA under Long Term Space Astrophysics Grant no. S-92654-F. Work at LLNL was performed under the auspices of the US Department of Energy, Contract mo. W-7405-Eng-48."
    },
    "0002/astro-ph0002168_arXiv.txt": {
        "abstract": "Gamma-ray burst astronomy has undergone a revolution in the last three years, spurred by the discovery of fading long-wavelength counterparts.  We now know that at least the long duration GRBs lie at cosmological distances with estimated electromagnetic energy release of $10^{51}$ -- $10^{53}$ erg, making these the brightest explosions in the Universe.  In this article we review the current observational state, beginning with the statistics of X-ray, optical, and radio afterglow detections.  We then discuss the insights these observations have given to the progenitor population, the energetics of the GRB events, and the physics of the afterglow emission.  We focus particular attention on the evidence linking GRBs to the explosion of massive stars.  Throughout, we identify remaining puzzles and uncertainties, and emphasize promising observational tools for addressing them.  The imminent launch of {\\em HETE-2} and the increasingly sophisticated and coordinated ground-based and space-based observations have primed this field for fantastic growth. ",
        "introduction": "\\label{sec:introduction}  GRBs have mystified and fascinated astronomers since their discovery. Their brilliance and their short time variability clearly suggest a compact object (black hole or neutron star) origin. Three decades of high-energy observations, culminating in the definitive measurements of CGRO/BATSE, determined the spatial distribution to be isotropic yet inhomogeneous, suggestive of an extragalactic population (see \\cite{fm95} for a review of the situation prior to the launch of the BeppoSAX mission).  Further progress had to await the availability of GRB positions adequate for identification of counterparts at other wavelengths.   In the cosmological scenario, GRBs would have energy releases comparable to that of supernovae (SNe). Based on this analogy, Paczy\\'nsk \\&\\ Rhoads \\cite{pr93} and Katz \\cite{k94} predicted that the gamma-ray burst would be followed by long-lived but fading emission. These papers motivated systematic searches for radio afterglow, including our effort at the VLA \\cite{fk95}. The broad-band nature of this ``afterglow'' and its detectability was underscored in later work \\cite{mr97,v97}.  Ultimately, the detection of the predicted afterglow had to await localizations provided by the Italian-Dutch satellite, BeppoSAX \\cite{b+97}.  The BeppoSAX Wide Field Camera (WFC) observes about 3\\% of the sky, triggering on the low-energy (2 -- 30 keV) portion of the GRB spectrum, localizing events to $\\sim$ 5 -- 10 arcminutes.  X-ray afterglow was first discovered by BeppoSAX in GRB~970228, after the satellite was re-oriented (within about 8 hours) to study the error circle of a WFC detection with the 2 -- 10 keV X-ray concentrators.  The detection of fading X-ray emission, combined with the high sensitivity and the ability of the concentrators to refine the position to the arcminute level, led to the subsequent discovery of long lived emission at lower frequencies \\cite{c+97,JvP+97,f+97}.  Optical spectroscopy of the afterglow of GRB 970508 led to the definitive demonstration of the extragalactic nature of this GRB \\cite{mdk+97}.  The precise positions provided by radio and/or optical afterglow observations have allowed for the identification of host galaxies, found in almost every case.  Not only has this provided further redshift determinations, but it has been useful in tying GRBs to star formation through measurements of the host star formation rate (e.g. \\cite{kdr+98,dkb+98}). HST with its exquisite resolution has been critical in localizing GRBs within their host galaxies and thereby shed light on their progenitors (e.g. \\cite{fpt+99,hh99,bod+99}). Observations of the radio afterglow have directly established the relativistic nature of the GRB explosions \\cite{f+97} and provided evidence linking GRBs to dusty star-forming regions. Radio observations are excellent probes of the circumburst medium and the current evidence suggests that the progenitors are massive stars with copious stellar winds.  The latest twist is an apparent connection of GRBs with SNe \\cite{bkd+99}. Separately, an important development is the possible association of a GRB with a nearby (40 Mpc) peculiar SN \\cite{gvv+98,kfw+98}.  \\begin{figure}[htb] \\centerline{\\psfig{figure=980703-lcurve.ps,width=7.3cm}\\qquad\\qquad\\psfig{figure=970508-lcurve.ps,width=7.3cm,angle=0}} \\vspace{10pt} \\caption[]{\\small {\\it Left: The radio light curve of GRB 980703. This is a typical afterglow, a rise to a peak followed by a power law decay.  The longer lifetime of the radio afterglow allows us to see both the rise and the fall of the afterglow emission. In contrast, at optical and X-ray emission, most of the times we see only the decaying portion of the light curve.  Right: The radio light curve of GRB970508 \\cite{fwk00}.  The wild fluctuations of the light curve in the first three weeks are chromatic.  At later times, the fluctuations become broad-band and subdued.  These fluctuations are a result of multi-path propagation of the radio waves in the Galactic interstellar medium. As the source expands (at superluminal speeds) the scintillation changes from diffractive to refractive scintillation. This is analogous to why stars twinkle but planets do not.\\label{fig:980703-970508}}} \\end{figure}  In this paper we review the primary advances resulting from afterglow studies. \\S{II} discusses the statistics of detections to-date, including possible causes for the lack of radio and optical afterglows from some GRBs. In \\S{III} we review constraints on the nature of the progenitor population(s), in particular evidence linking some classes of GRBs to SNe.  \\S{IV} describes the status of current understanding of the physics of the afterglow emission.  Here we compare observations to predictions of the basic spherically-symmetric model, and describe complications arising from deviations from spherical symmetry and non-uniform distribution of the circumburst medium.  We conclude with speculations of the near and long-term advances in this field (\\S{V}).  We point out that this review has two biases. First, given the concentration of previous review articles on optical and X-ray observations, we emphasize the unique contributions of radio afterglow measurements. Second, this article is intended to also provide a summary of the efforts of the Caltech-NRAO-CARA GRB collaboration, and therefore details our work in particular. This review is in response to review talks given at the 1999 Maryland October meeting (SRK) and the 5th Huntsville GRB meeting (DAF and SRK). ",
        "conclusions": ""
    },
    "0002/astro-ph0002442_arXiv.txt": {
        "abstract": "In magnetic fields stronger than $B_Q \\equiv m_e^2 c^3/\\hbar e = 4.4 \\times 10^{13}$ Gauss, an electron's Landau excitation energy exceeds its rest energy. I review the physics of this strange regime and some of its implications for the crusts and magnetospheres of neutron stars. In particular, I describe how ultra-strong fields \\begin{list}{$\\bullet$}{\\setlength{\\leftmargin}{1.8\\leftmargini} \\setlength{\\rightmargin}{\\leftmargin} \\setlength{\\itemsep}{0pt} \\setlength{\\parsep}{0pt} \\setlength{\\topsep}{2pt} } \\item render the vacuum {\\it birefringent\\/} and capable of distorting and magnifying images (``magnetic lensing\"); \\item change the self-energy of electrons: as $B$ increases they are first slightly lighter than $m_e$, then slightly heavier; \\item cause photons to rapidly split and merge with each other; \\item distort atoms into long, thin cylinders and molecules into strong, polymer-like chains; \\item enhance the pair density in thermal pair-photon gases; \\item strongly suppress photon-electron scattering, and \\item drive the vacuum itself unstable, at extremely large $B$. \\end{list} In a concluding section, I discuss the spindown of ultra-magnetized neutron stars and recent soft gamma repeater observations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002397_arXiv.txt": {
        "abstract": "We  evaluate the observational constraints on the  spectral index $n$, in the context of the $\\Lambda$CDM hypothesis which represents the simplest viable cosmology. We first take $n$ to be  practically scale-independent. Ignoring reionization, we find at  a nominal 2-$\\sigma$ level $n\\simeq 1.0 \\pm 0.1$. If we make the more realisitic assumption that reionization occurs when a  fraction $f\\sim 10^{-5}$ to $1$ of the matter has collapsed, the 2-$\\sigma$  lower bound is unchanged while the 1-$\\sigma$ bound rises slightly. These constraints  are compared  with the prediction of various inflation models. Then we investigate the two-parameter scale-dependent spectral index, predicted by  running-mass inflation models, and find that present data allow significant scale-dependence of $n$, which occurs in a physically reasonable regime of parameter space. ",
        "introduction": "It is generally supposed that structure in the Universe originates from a primordial gaussian curvature perturbation,  generated by slow-roll inflation. The  spectrum $\\calpr(k)$ of the curvature perturbation is the point of contact between observation and models of inflation. It is   given in terms of the inflaton potential $V(\\phi)$ by\\footnote {As usual, $\\mpl=2.4\\times 10^{18}\\GeV$ is the Planck mass, $a$ is the scale factor and $H=\\dot a/a$ is the Hubble parameter, and $k/a$ is the wavenumber. We assume the usual slow-roll conditions $\\mpl^2|V''/V|\\ll 1$ and $\\mpl^2(V'/V)^2\\ll1$, leading to $3H\\dot\\phi\\simeq -V'$.} \\be \\frac4{25}\\calpr(k) = \\frac1{75\\pi^2\\mpl^6}\\frac{V^3}{V'^2} \\,, \\label{delh} \\ee where the potential and its derivatives are evaluated at the epoch of horizon exit $k=aH$. To work out the value of $\\phi$ at this epoch one uses the relation \\be \\ln(k\\sub{end}/k)\\equiv N(k) =\\mpl^{-2}\\int^\\phi_{\\phi\\sub{end}} (V/V') \\diff\\phi \\,, \\label{Nofv} \\ee where $N(k)$ is actually the number of $e$-folds from horizon exit  to the  end of slow-roll inflation. At the scale explored by the COBE measurement of the cosmic microwave background (cmb) anisotropy, $N(k\\sub{COBE})$ depends on the expansion of the Universe after inflation in the manner specified by \\eq{Ncobe} below.  Given this prediction, the observed large-scale  normalization  $\\calp_\\calr^{1/2}\\simeq 10^{-5}$ provides a strong constraint on models of inflation. Taking that for granted, we are here interested in the scale-dependence of the spectrum, defined by the, in general, scale-dependent spectral index $n$; \\be n(k)-1\\equiv \\frac {\\diff \\ln \\calpr}{ \\diff \\ln k} \\,. \\ee According to most inflation models, $n$ has negligible variation on cosmological scales so that $\\calpr\\propto k^{n-1}$, but we shall also discuss an interesting class of models giving a different scale-dependence.  {}From \\eqs{delh}{Nofv}, \\bea n-1 &=&  2\\mpl^2 (V''/V)-3\\mpl^2 (V'/V)^2 \\,, \\label{nofv} \\eea and in  almost all models of inflation, \\eq{nofv} is well approximated by \\be n-1=2\\mpl^2(V''/V) \\label{nofvapprox} \\,. \\ee We see that the  spectral index measures the {\\em shape} of the inflaton potential $V(\\phi)$, being independent of its overall normalization. For this reason, it is a powerful discriminator between models of inflation.  The observational constraints on the spectral index have been studied by many authors, but a new investigation is justified for  two reasons. On the observational side, the cosmological parameters are at last being pinned down, as is  the height of the first peak in the spectrum the cmb anisotropy. No study has yet been given which takes on board  these observational developments, while at the same time taking on board the crucial influence of the reionization epoch on the peak height. On the theory side, it is known  that the spectral index may  be strongly scale-dependent if  the inflaton has a gauge coupling, leading to what are called running-mass models. The quite specific, two-parameter prediction for the scale dependence of the  spectral index in these models has not been compared with presently available data. ",
        "conclusions": "In the context of the $\\Lambda$CDM model, we have evaluated the observational constraint on the  spectral index $n(k)$. This constraint comes from a range of data, including the height of the first peak in the cmb anisotropy,  which we take to be $80\\pm 10\\mu$K (nominal 1-$\\sigma$). Reionization is assumed to occur when some fixed fraction $f$ of the matter collapses, and the most important results are insensitive to this fraction in the reasonable range $10^{-4}\\lsim f\\lsim 1$.  We first considered the  case that $n$ has negligible scale dependence, comparing the observational bound with the prediction of various models of inflation. A significant improvement in the 2-$\\sigma$ lower bound, which may well occur with the advent of slightly better measurements of the cmb anisotropy, will   become a serious discriminator between models of inflation. Even the present  bound has serious implications if, as is very possible,  late-time gravitino creation or some other phenomenon requires an era of thermal inflation after the usual inflation.  We also considered  the running mass models of inflation, where the spectral index can have significant scale-dependence. Because of this scale dependence, it is in this case crucial to fix not the epoch of reionization, but the fraction $f$ of matter that has collapsed at that epoch. We presented results for the choice $f=1$ (corresponding to $z\\sub R\\simeq 13$ if the spectral index has negligible scale-dependence), and for a perhaps more reasonable choice  $ f=10^{-2.2}$. In the running-mass  models, the scale-dependent spectral index  $n(k)$ is given by $n-1=s \\exp(c\\Delta N) -c$, where $\\Delta N=\\ln(k\\sub{COBE}/k)$. The parameters in this expression can be of either sign, leading to four different models of inflation. Barring fine-tuning, one expects  $s$ to be in the range $|c|e^{-cN\\sub{COBE}}\\lsim |s|\\lsim e^{-cN\\sub{COBE}}$. The parameter $c$ depends on the nature of the soft supersymmetry breaking, but in the simplest case of gravity-mediation it becomes a dimensionless coupling strength, presumably of order $ 10^{-1}$ to $10^{-2}$ in magnitude.  Without worrying  about the origin of the parameters $c$ and $s$, we have investigated the observational constraints  on them. In  the  case $c,s>0$ (referred to as Model (i)) we  find that $n$ can have a significant  variation on cosmological scales, with $n-1$ passing  through zero signaling a  minimum of the spectrum of the primordial curvature perturbation. In a future paper, we shall exhibit  the possible  effect of this scale-dependence on the cmb anisotropy, at and above the first peak."
    },
    "0002/astro-ph0002504_arXiv.txt": {
        "abstract": "The clustering of galaxies is well characterized by fractal properties, with the presence of an eventual cross-over to homogeneity still a matter of considerable debate. In this letter we discuss the cosmological implications of a fractal distribution of matter, with a possible cross-over to homogeneity at an undetermined scale $R_{homo}$. Contrary to what is generally assumed, we show that, even when $R_{homo} \\rightarrow \\infty$, this possibility can be treated consistently within the framework of the expanding universe solutions of Friedmann. The fractal is a perturbation to an open cosmology in which the leading homogeneous component is the cosmic background radiation (CBR). This cosmology, inspired by the observed galaxy distributions, provides a simple explanation for the recent data which indicate the absence of deceleration in the expansion ($q_o \\approx 0$). Correspondingly the `age problem' is also resolved. Further we show that the model can be extended back from the curvature dominated arbitrarily deep into the radiation dominated era, and we discuss qualitatively the modifications to the physics of the anisotropy of the CBR, nucleosynthesis and structure formation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002218_arXiv.txt": {
        "abstract": "The time profiles of many gamma-ray bursts consist of distinct pulses, which offers the possibility of characterizing the temporal structure of these bursts using a relatively small set of pulse shape parameters.  We have used a pulse decomposition procedure to analyze the Time-to-Spill (TTS) data for all bursts observed by BATSE up through trigger number 2000, in all energy channels for which TTS data is available.  We obtain amplitude, rise and decay timescales, a pulse shape parameter, and the fluences of individual pulses in all of the bursts.  We investigate the correlations between brightness measures (amplitude and fluence) and timescale measures (pulse width and separation) which may result from cosmological time dilation of bursts, or from intrinsic properties of burst sources or from selection effects.  The effects of selection biases are evaluated through simulations.  The correlations between these parameters among pulses within individual bursts give a measure of the intrinsic effects while the correlations among bursts could result both from intrinsic and cosmological effects.  We find that timescales tend to be shorter in bursts with higher peak fluxes, as expected from cosmological time dilation effects, but also find that there are non-cosmological effects contributing to this inverse correlation.  We find that timescales tend to be longer in bursts with higher total fluences, contrary to what is expected from cosmological effects.  We also find that peak fluxes and total fluences of bursts are uncorrelated, indicating that they cannot both be good distance indicators for bursts. ",
        "introduction": "Many of the signatures of the cosmological time dilation and the radiation mechanisms of gamma-ray bursts (GRBs) are hidden in the temporal and spectral characteristics of GRBs.  The subject of this paper is the analysis of the temporal properties of the bursts, and the correlations between intensities and timescales.  We use the BATSE Time-to-Spill (TTS) data, which can give much higher time resolution than other forms of BATSE data for most bursts.  The advantages and shortcomings of this data, our decomposition of the time profiles into pulses, and the evolution of burst characteristics are described in greater detail in the accompanying paper \\cite{lee:2000}.  What follows is a brief summary.  (See also \\cite{lee:1996,lee:1998,lee:thesis}.)  Many burst time profiles appear to be composed of a series of discrete, often overlapping, pulses, often with a \\emph{fast rise, exponential decay} (FRED) shape~\\citep{norris:1996}.  The different pulses may represent emission from distinct subevents within the gamma-ray burst source.  Therefore, it may be useful to decompose burst time profiles in terms of individual pulses, each of which rises from background to a maximum and then decays back to background levels.  We have analyzed gamma-ray burst time profiles by representing them in terms of a finite number of pulses, each of which is described by a small number of parameters.  We have used the phenomological pulse model of \\cite{norris:1996} to decompose gamma-ray burst time profiles into distinct pulses.  In this model, each pulse is described by five parameters with the functional form \\begin{equation} I(t)\\  =\\  A \\exp{\\left(-\\left\\vert\\frac{t - t_{\\text{max}}}{\\sigma_{r,d}}\\right\\vert^{\\nu}\\right)}\\ , \\label{eq:pulse} \\end{equation} where $t_{\\text{max}}$ is the time at which the pulse attains its maximum, $\\sigma_{r}$ and $\\sigma_{d}$ are the rise and decay times, respectively, $A$ is the pulse amplitude, and $\\nu$ (the~``peakedness'') gives the sharpness or smoothness of the pulse at its peak.  We have developed an interactive pulse-fitting program to perform this pulse decomposition on the BATSE TTS data. and used this program to fit pulses to all gamma-ray bursts in the BATSE 3B catalog~\\citep{batse:3b} up to trigger number 2000 in all of the four BATSE LAD energy channels for which TTS data is available and shows time variation beyond the normal Poisson noise for the background.  We fit each channel of each burst separately.  We have obtained 574 fits for 211 bursts, with a total of 2465 pulses.  In this paper, we focus on the possibility of distinguishing between intrinsic signatures in the temporal characteristics and those which arise from their cosmological distribution.  A prominent example of this is the cosmological time dilation effect, which we expect to see since some, and possibly all, gamma-ray bursts originate at cosmological distances.  All timescales in GRBs will be lengthened by a factor of $1 + z$ where $z$ is the redshift of the burst, as a result of cosmological time dilation~\\citep{paczynski:1992,piran:1992}.  However, this seemingly straightforward test is not simple.  First of all, given the great diversity in burst time profiles, it is difficult to decide which timescale is most appropriate for this test.  It seems unlikely that any particular timescale is approximately the same in all bursts, so we expect to find time dilation as a statistical effect, rather than for individual bursts.  Secondly, redshifts are known only for a few bursts, so that for the vast majority of bursts we need to use another measure of distance or redshift.  Most past analyses have used some measure of apparent GRB brightness for this purpose with the tacit assumption that the corresponding intrinsic brightness is a standard candle or has a very narrow distribution.  The observed apparent brightnesses of bursts are generally measured using either peak fluxes, which give the instantaneous intensity of bursts when they peak, or fluences, which measure the total output of bursts integrated over their entire durations.  The brightness measures can also be divided another way, into photon measures and energy measures.  Thus, there are several different measures of the apparent brightnesses of bursts.  The BATSE burst catalogs give peak photon fluxes and total energy fluences for bursts.  The pulse-fitting data presented here can be used to determine count fluxes and count fluences.  Most previous work on the evidence for time dilation in burst time profiles has binned the bursts into two or three brightness classes using the peak flux as a measure of brightness, and compared a measure of total burst duration these classes.  Use of fluence as a brightness measure has been promoted by \\cite{petrosian:1996} and \\cite{lloyd:1999}.  In this paper, we use a number of different timescale and brightness measures.  We will describe their correlations using power laws. Although cosmological models generally predict more complex relationships than a simple power law, it would be fruitless to attempt to fit anything more complex than a power law using the pulse-fitting data, which appears to have a large intrinsic scatter. To contrast the cosmological versus the intrinsic signatures, we compare the relations or correlations between strengths and timescales among bursts, which should contain the signatures of cosmological time dilations, with the same correlations among pulses of individual bursts, which can only contain the intrinsic effects.  It is likely that some of these correlations are affected by selection effects in our fitting procedures.  To investigate the importance of these, we have carried out extensive simulations which are described in the accompanying paper \\cite{lee:2000}.  We use the results of these simulations to test whether or not the correlations we find are properties of the bursts or are products of our procedures.  In the next section, we define the various timescales and burst strengths used in this analysis.  The correlations relevant to the ``time dilation'' tests are discussed in Section~\\ref{sec:timedilation} and the correlations between other quantities within bursts and among bursts are described in Section~\\ref{sec:othercorr}.  In Section~\\ref{sec:discuss} we discuss the significance of these correlations.  It should be noted that many of the simulated bursts were affected by a truncation that almost never occurred in the actual BATSE TTS data. The TTS data is truncated at $2^{20}$ counts or 240 seconds, whichever occurs first.  In nearly all of the actual bursts, the 240 second limit is reached first, while in many of the the simulated bursts, the $2^{20}$ count limit is reached first.  This truncation can shorten the observed time intervals between the first and last pulses in a burst, and between the two highest amplitude pulses in a burst, but not the observed pulse widths or the observed time intervals between consecutive pulses.  \\emph{Therefore, all discussions of the first two kinds of time intervals in simulated bursts will only consider simulated bursts where no pulses were truncated by the $2^{20}$ count limit.} ",
        "conclusions": "\\label{sec:discuss}  In this paper, we use a pulse-fitting procedure to the TTS data from BATSE and determine the amplitudes, rise and decay times, and fluences.  We investigate the correlations between all of these parameters of pulses in individual bursts and among different bursts. The former gives a measure of correlations intrinsic to the energy and radiation generation in burst sources, while the latter are also affected by cosmological effects.  Simulations are used to determine the biases of the pulse-fitting procedure.  If the peak luminosities of pulses or bursts are approximate standard candles, so that the peak fluxes would be good measures of distance, then we expect to find negative correlations between fluxes and timescales.  We do find inverse correlations between the highest pulse amplitude within a burst and two different timescales, the width of the highest amplitude pulse and the time interval between the two highest amplitude pulses.  The former correlation, between pulse amplitude and pulse width, which is expected from cosmological time dilation effects, is nevertheless not consistent with purely cosmological effects, but must be at least partially influenced by non-cosmological effects.  These non-cosmological effects may include intrinsic properties of the burst sources, or selection effects due to the BATSE triggering procedure, but do not appear to be affected by the pulse-fitting procedure.  Our study indicates that the latter correlation, between pulse amplitude and time intervals between pulses, may be less influenced by non-cosmological effects.  The inverse correlation observed between pulse amplitude and pulse width within bursts results in part from selection effects in the pulse-fitting procedure, but also appears to result in part from intrinsic properties of the burst sources.  If the total radiated energies of bursts are approximate standard candles, so that the burst fluences would be good measures of distance, then we expect to find negative correlations between fluences and timescales.  We find instead a \\emph{positive} correlation between the total burst count fluence and the width of the highest amplitude pulse, but no correlation with the time interval between the two highest amplitude pulses.  The former correlation indicates that non-cosmological effects are stronger than any cosmological effects.  This is supported by the positive correlation between pulse amplitude and pulse count fluence within bursts. However, it is not clear why total burst count fluence and time intervals between pulses show no correlation.  It is natural to expect that the peak flux of bursts and the total count fluence of bursts should both decrease essentially the same way (except for a factor of $1 + z$) as the distance to the burst sources increase.  This would suggest that there should be positive correlations between the peak flux of bursts and the total count fluence of bursts.  Strangely, the highest pulse amplitude and the total count fluence of bursts appear to have no statistically significant correlation with each other, implying that the two measures of brightness cannot both be good standard candles; at least one, or more probably both, are poor measures of distance.  There do not appear to be any statistically significant correlations between pulse amplitude and pulse asymmetry, whether the comparison is i) of all pulses in all bursts combined, ii) of only the highest pulse in each burst, iii) of only the single-pulse bursts, or iv) of different pulses within multiple-pulse bursts.  This implies that the differences between the variations of pulse rise and decay time with pulse amplitude are statistically insignificant, and both rise times and decay times tend to decrease as pulse amplitude increases."
    },
    "0002/hep-ph0002031_arXiv.txt": {
        "abstract": "s{Cosmology cannot rule out the solution $\\nu_{\\mu}\\leftrightarrow\\nu_{s}$ to the atmospheric neutrino data and thus only Earth experiments will be able to give a definitive answer. This conclusion holds when a generation of lepton number is taken into account and one assumes that the sterile neutrino is also slightly mixed with an ${\\rm eV}-\\tau$ neutrino. This result cannot be spoiled by a chaotic generation of lepton domains.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002510_arXiv.txt": {
        "abstract": "We present the microlensing optical depth towards the Galactic bulge based on the detection of 99 events found in our Difference Image Analysis (DIA) survey. This analysis encompasses three years of data, covering $\\sim 17$ million stars in $\\sim 4$ deg$^2$, to a source star baseline magnitude limit of $V = 23$. The DIA technique improves the quality of photometry in crowded fields, and allows us to detect more microlensing events with faint source stars. We find this method increases the number of detection events by $85\\%$ compared with the standard analysis technique. DIA light curves of the events are presented and the microlensing fit parameters are given. The total microlensing optical depth is estimated to be $\\tau_{total}= 2.43^{+0.39}_{-0.38}\\times 10^{-6}$ averaged over 8 fields centered at $l=2\\fdg68$ and $b=-3\\fdg35$. % For the bulge component we find $\\tau_{bulge}=3.23^{+0.52}_{-0.50}\\times 10^{-6}$ assuming a $25\\%$ stellar contribution from disk sources.  These optical depths are in good agreement with the past determinations of the MACHO \\shortcite{ALC97a} and OGLE \\shortcite{USKK94d} groups, and are higher than predicted by contemporary Galactic models.  We show that our observed event timescale distribution is consistent with the distribution expected from normal mass stars, if we adopt the stellar mass function of Scalo (1986) as our lens mass function. However, we note that as there is still disagreement about the exact form of the stellar mass function, there is uncertainty in this conclusion.  Based on our event timescale distribution we find no evidence for the existence of a large population of brown dwarfs in the direction of the Galactic bulge. ",
        "introduction": "Over the past seven years the MACHO group has been making observations of the Galactic bulge in order to determine some of the fundamental properties of our Galaxy. The Milky Way is expected to be an SAB(rs)bc or SAB(r)bc type spiral galaxy \\shortcite{DEV64,Fux97} with four spiral arms \\shortcite{VA95}. However, very little is known about the mass distributions of the various components of our Galaxy (bulge, spheroid, disk, halo).  Galactic microlensing surveys provide some insight into the structure and dynamics of the inner Galaxy, spiral arms and the halo. Unlike most types of observation, the presence of lensing objects can be detected independent of their luminosities.  Microlensing is sensitive to the mass distribution rather than light, this makes microlensing a powerful way of investigating the mass density within our Galaxy. Furthermore, microlensing can be used to investigate the stellar mass function to the hydrogen burning limit, both within our Galaxy and other nearby galaxies.  The amplification of a source star during gravitational microlensing is related to the projected lens-source separation $u$ normalised by the angular Einstein Ring radius $R_{E}$. This is given by  \\begin{equation}\\label{chisw} A(u) = \\frac{u^{2}+2}{u \\sqrt{u^{2}+4}}. \\end{equation}  \\noindent The timescale of a microlensing event, $\\hat t$, is characterised by the time it takes for the Einstein ring associated with a foreground compact lensing object, to transit a background source star at velocity $v_{\\perp}$.  The size of the Einstein ring, for a lens with mass $M$ (in $\\rm M_{\\odot}$), an observer-lens distance $D_{d}$, and a source-observer distance $D_{s}$, is given by  \\begin{equation} R_{E} = 2.85 {\\rm AU} \\sqrt{\\frac{M D_{d}(1-\\frac{D_{d}}{D_{s}})}{1\\,\\rm kpc}}. \\label{chiswolsonx} \\end{equation}  \\noindent The lensing timescale is $\\hat t \\equiv 2R_{E}/v_{\\perp}$. Hence, if $R_{E}$ were known, this would enable us to constrain some of the physical parameters of a microlensing event ($M$, $D_{s}$, $D_{d}$). However, $R_{E}$ is not known, generally, so it is not possible to determine the lens masses from individual microlensing events.  Nevertheless, under special circumstances it is possible to impose additional constraints on these microlensing event parameters when quantities such as, the physical size of the source star \\cite{AAAA99} or the projected transverse velocity of the lens \\cite{ALC95a} are measured.  Photometry of the stars monitored by the MACHO project has previously been carried out using a fixed position PSF photometry package SoDoPhot (Son of DoPhot, \\shortciteNP{BEN93}).  In 1996 we introduced a second reduction method, Difference Image Analysis (hereafter DIA). The DIA technique enables us to detect microlensing events which go undetected with the SoDoPhot photometry because the events are due to stars which are too faint to be detected when unlensed. This technique follows on from the work of \\citeN{Crotts92}, \\citeN{PHL95} and \\citeN{TC96}, and allows us to detect and perform accurate photometry on these new microlensing events found in the reanalysis of bulge images.   Recently, the MACHO and OGLE groups reported that the microlensing optical depth towards the Galactic bulge was a factor of 2 larger than expected from stellar number density.  That is, the optical depth found by OGLE is $3.3\\pm{1.2} \\times 10^{-6}$ \\shortcite{USKK94d} and by MACHO is $3.9^{+1.8}_{-1.2} \\times 10^{-6}$ for 13 clump giant source star events out of a 41 event sample \\shortcite{ALC97a}.  It was suggested that the size of these measurements could be explained by the presence of a bar oriented along our line-of-sight to the bulge (\\shortciteNP{PSUS94b,ZSR95}). The density profile of the proposed bar is given by  \\begin{equation}\\label{pt1} \\rho_{b} = \\frac{M}{20.65 abc}\\; {\\rm exp}\\left (-\\frac{w^{2}}{2}\\right), \\end{equation}  \\noindent where  \\begin{equation}\\label{pt2} w^{4} = \\left[\\left(\\frac{x}{a}\\right)^{2} + \\left(\\frac{y}{b}\\right)^{2}\\right]^{2} + \\left(\\frac{z^{\\prime}}{c}\\right)^{4}. \\end{equation}  \\noindent For the bulge galactocentric coordinates (x,y,z$^{\\prime}$): $x = \\cos{\\theta} - \\eta \\cos{b} \\cos{(l-\\theta)}$, $y=  -\\sin{\\theta} - \\eta \\cos{b}\\sin{(l-\\theta)}$, $z^{\\prime} = \\eta \\sin{b}$. The bar inclination angle $\\theta$ is oriented in the direction of increasing $l$, and $\\eta=D_{s}/D_{8.5}$ is the ratio of the source distance relative to a galactocentric distance, taken to be 8.5 kpc. The terms $a$, $b$ and $c$ define the bar scale lengths.  The idea that our Galaxy harbours a bar at its centre is not a new one as it was first suggested by \\shortciteN{DEV64} because of the similarity of the gas dynamics observed in our galaxy with other barred galaxies. \\citeN{BGSBU91} provided further evidence for a bar from star counts. The DIRBE results of \\citeN{DAH95} were also found to be consistent with this prediction. The presence of such a bar is an important way of explaining the interaction of the disk, halo and the spiral density waves in the disk.  A number of authors have adopted a bar into their Galactic models and have adopted various values of the bar orientation \\shortcite{PSUS94b,Peale98,ZM96} and bar mass \\shortcite{Peale98,ZM96}. Other authors have also proposed that large optical depth contributions could come from the disk component \\shortcite{EGTB98}, or the Galactic stellar mass function (M\\'era, Chabrier, \\& Schaeffer 1998; \\shortciteNP{HC98,ZM96}).\\nocite{MER98a}  In this paper we present a new value for the microlensing optical depth and investigate what is known about the Galactic parameters with most influence the optical depth. In the next section we will detail the observational setup. In \\S 3, we will outline the reduction procedure.  We will next review the microlensing event selection process in \\S 4.  The results of our analysis are presented in \\S 5, and we will discuss how the microlensing detection efficiency was calculated in \\S 6.  The microlensing optical depth for the sample of fields presently analysed with DIA and for each of the individual field will be presented in \\S 7.  In \\S 8, we will review what is known about the most important factors affecting the observed optical depth and discuss the implications of our results.  In the final section we summarise the results of this work. ",
        "conclusions": ""
    },
    "0002/astro-ph0002383_arXiv.txt": {
        "abstract": "Future space and ground-based survey programmes will produce an impressive amount of photometric data. The GAIA space mission will map the complete sky down to mag V=20 and produce time series for about 1 billion stars. Survey instruments as the International Liquid Mirror Telescope will observe slices of the sky down to magnitude V=23. In both experiments, the opportunity exists to discover a huge amount of variable stars. A prediction of the expected total number of variable stars and the number of variables in specific subgroups is given. ",
        "introduction": "A first estimate of the total number of variable stars observable by GAIA was done by Eyer (1999). The star population used came from the star-count model of Figueras~et~al.~(1999) and the variability detection threshold was derived from the Hipparcos survey results. With the new qualifications of the GAIA mission, about 1 billion stars (up to mag~G$<$20) are expected to be observed, with about 18~million variable stars, including about 5~million \"classic\" periodic variables.\\\\ Very different star counts are obtained according to the extinction laws used (Figueras, private communication). Since the quality of the GAIA photometry in the crowded fields is still uncertain, we cannot discuss here the number of variables in dense clusters and galaxies.\\\\ About 2 to 3 millions eclipsing binaries will be observed, but their detection probability will be studied in detail in the future. About 300\\,000 stars with rotation induced variability can be expected as well. ",
        "conclusions": ""
    },
    "0002/astro-ph0002456_arXiv.txt": {
        "abstract": "Using the Short-Wavelength-Spectrometer on the Infrared Space Observatory (ISO), we obtained near- and mid-infrared spectra toward the brightest H$_2$ emission peak of the Orion OMC-1 outflow. A wealth of emission and absorption features were detected, dominated by 56 $\\rm H_2$ ro-vibrational and pure rotational lines reaching from H$_2$ 0--0 S(1) to 0-0~S(25).  The spectra also show a number of H~{\\sc i} recombination lines, atomic and ionic fine structure lines, and molecular lines of CO and $\\rm H_2O$.  Between 6 and 12 $\\rm \\mu m$ the emission is dominated by PAH features.  The extinction toward the molecular and atomic line emitting regions is estimated from relative line intensities, and it is found that the H$_2$ emission arises from within the OMC-1 cloud at an average K-band extinction of 1.0 mag, whereas the atomic hydrogen emission and much of the fine structure emission comes from the foreground H~{\\sc ii} region and its bounding photodissociation front.  The total H$_2$ luminosity in the ISO-SWS aperture is estimated at $(17 \\pm 5)~\\rm L_{\\sun}$, and extrapolated to the entire outflow, $(120 \\pm 60)~\\rm L_{\\sun}$.  The H$_2$ level column density distribution shows no signs of fluorescent excitation or a deviation from an ortho-to-para ratio of three.  It shows an excitation temperature which increases from about 600~K for the lowest rotational and vibrational levels to about 3200~K at level energies $E(v,J)/k > 14\\,000$~K. No single steady state shock model can reproduce the observed H$_2$ excitation.  The higher energy H$_2$ levels may be excited either thermally in non-dissociative J-shocks, through non-thermal collisions between fast ions and molecules with H$_2$ in C-shocks, or they are pumped by newly formed H$_2$ molecules.  The highest rotational levels may be populated by yet another mechanism, such as the gas phase formation of H$_2$ through H$^-$. ",
        "introduction": "The Orion molecular cloud, OMC-1, located behind the Orion M42 Nebula at a distance of $\\sim$450~pc (Genzel \\& Stutzki \\cite{gen89}), is the best-studied massive star forming region.  This cloud embeds a spectacular outflow arising from some embedded young stellar object, which can possibly be identified as the radio source ``I'' 0.49 arcsec south of the infrared source IRc2-A (Menten \\& Reid \\cite{men95}; Dougados et al. \\cite{dou93}).  The outflow shocks the surrounding molecular gas, thereby giving rise to the strongest H$_2$ infrared line emission appearing in the sky (Fig.~\\ref{schultz}).  Peak~1 (Beckwith et al. \\cite{bec78}) is the brighter of the two H$_2$ emission lobes of the outflow.  Although the outflow has been studied extensively for nearly two decades, the nature of the excitation mechanism remains unclear.  Molecular hydrogen, through its infrared rotational and rotation-vibrational transitions, is an important coolant in shocks and photodissociation regions, and thereby a particularly well suited tracer of the flourescently- and/or shock-excited gas.  The Short Wavelength Spectrometer (SWS, de Graauw et al. \\cite{deg96}) aboard the Infrared Space Observatory (ISO, Kessler et al. \\cite{kes96}) offered the first opportunity to observe pure rotational and rotation-vibrational H$_2$ lines from 2.4 $\\mu$m to 28 $\\mu$m with one instrument, unhindered by the Earth's atmosphere.  In this paper, we present a comprehensive set of intensities for 56 H$_2$ near- and mid-infrared lines we observed with the ISO-SWS.  These observations trace populations of energy levels ranging from $E/k=$1015~K to 43\\,000~K.  The redundancy of the H$_2$ level determinations provides information on the average gas excitation along the line of sight over an unprecedented range.  This sheds new light on the possible excitation mechanisms in the IRc2 outflow.  We here concentrate on the interpretation of the $\\rm H_2$ and the atomic and ionic fine structure line emission, whereas a detailed discussion of the CO and $\\rm H_2O$ lines will be presented in a separate paper (Boonman et al.  \\cite{boo00}).  In a related paper we already discussed the detection of HD toward Orion Peak~1 (Bertoldi et al. \\cite{ber99}). ",
        "conclusions": "\\begin{figure*}[htb] \\begin{center} \\includegraphics[width=2.\\columnwidth]{figure2.ps} \\caption{2.4--45 \\mum\\ spectrum of Orion Peak~1 obtained in the SWS~01 grating scan observing mode. Some of the detected lines, bands and features are identified. The continuum levels of the individual bands, which differ due to aperture changes, were adjusted to make the spectrum appear continuous.} \\label{spectrum} \\end{center} \\end{figure*}  Fig.~\\ref{spectrum} shows the full SWS~01 spectrum.  Most of the observed continuum flux is probably coming from the strong Becklin Neugebauer (BN) source near the edge of the aperture. Aperture size changes from one band to another then cause changes in the intercepted continuum which are not simply proportional to the aperture size. In addition, there is extended continuum emission all over the outflow.  We normalized the line and continuum fluxes by the aperture size, assuming that there is uniform surface brightness at least for the line emission. The exact aperture profiles for the various wavelength bands is yet to be determined.  Assuming an effective aperture resembling those shown in Fig.~\\ref{schultz} is approximate. The error from this assumption, and the nonuniform continuum surface brightness cause additional relative offsets in the continuum fluxes of neighboring bands of --10\\% for the 7--12~$\\mu$m band, $-30\\%$ for the 12--16~$\\mu$m band, +15\\% at 16--19.5~$\\mu$m, -5\\% at 19.5--27.5~$\\mu$m, -7.5\\% at 27.5--29.5~$\\mu$m, and +5\\% above 29.5~$\\mu$m.  Since the observed H$_2$ line intensities result in a smooth distribution of column densities in the excitation diagram, Fig. \\ref{excit}, the line intensities appear not to be affected much by the uncertainty of the aperture.  Fig.~\\ref{spec} shows the SWS~01 spectrum in more detail.  Fig.~\\ref{multi} shows selected lines at higher spectral resolution from line scan observation in the SWS~02 and SWS~07/SWS~06 modes.  The SWS~06 grating spectra were simultaneously recorded with the SWS~07 Fabry-Perot spectra.  \\begin{figure*}[h] \\begin{center} \\includegraphics[width=1.95\\columnwidth]{figure3.ps} \\vspace{0.2cm} \\caption{The full scan SWS 01 spectrum of Fig. 2 in detail.} \\label{spec} \\end{center} \\end{figure*}  \\begin{figure*}[h] \\begin{center} \\includegraphics[width=2.0\\columnwidth]{figure4.ps} \\vspace{0.5cm} \\caption{Line scans in the SWS 02 or SWS 07/SWS 06 modes.  In cases where the flux from the SWS 01 full scan spectrum differs much from that found in other observing modes, or where lines do not show well in Figs.~\\ref{spectrum} and \\ref{spec}, we overplot the SWS 01 spectrum for comparison.} \\label{multi} \\end{center} \\end{figure*} The Peak 1 spectrum is dominated by a large number of rotational and ro-vibrational $\\rm H_2$ lines.  The pure rotational lines arise from levels with energies ranging from $E/k=$1015~K for the 0-0~S(1) line to $E/k=42\\,515$~K for the 0-0~S(25) line. They represent gas with excitation temperatures ranging from 600~K for the low energy levels to over 3000~K for level energies $E/k > 14\\,000$~K. The observed fluxes of the identified H$_2$ lines are listed in Table~\\ref{h2_table}.   The spectrum is rich also in H~{\\sc i} recombination lines and atomic and ionic fine structure lines.  Between 4.5 to 5~$\\mu$m, we find a forest of gaseous 1-0 ro-vibrational CO emission, possibly mixed with absorption of solid CO (van Dishoeck et al. \\cite{dis98}). Gaseous water is seen in emission through the $\\nu_2$ bending mode between 6.3 and 7~$\\mu$m and several lines between 30 and 45~$\\mu$m, and gaseous $\\rm CO_2$ is detected at 15~$\\mu$m. PAH features dominate the emission between 6 and 12~$\\rm \\mu m$.  Absorption features of water ice are seen at 3.1~$\\mu$m, of $\\rm CO_2$ ice at 4.25~$\\rm \\mu m$, and of silicate at 9.7~$\\rm \\mu m$.  The various observed lines and features probe different regions -- both within the SWS aperture and along the line of sight. The H~{\\sc ii} region in the foreground of OMC-1 contributes to the H recombination and ionic fine structure emission, whereas the PAH (Verstraete et al. \\cite{ver96}) emission and a large fraction of the [Si~{\\sc ii}]34.8$\\rm \\mu m$ emission (Haas et al. \\cite{haa91}) originate in the PDR between the H~{\\sc ii} region and the molecular cloud which embeds OMC-1. The PDR also contributes to the $\\rm H_2$ emission.  The shocks dominate the emission of H$_2$, H$_2$O and CO, and may make a minor contribution to the H recombination and most fine structure emission.   \\subsection{Observed H$_2$ level column densities}  All molecular hydrogen lines, due to the small radiative transition probabilities, remain optically thin.  Therefore the corresponding ``observed'' upper level column density can be computed from the observed line flux, \\begin{equation} \\label{eq:col} N_{\\rm obs}(v,J) = {4\\pi \\lambda \\over h c} {I_{\\rm obs}(v,J\\rightarrow v',J')\\over A(v,J\\rightarrow v',J')}~, \\end{equation} where $I_{\\rm obs}(v,J\\rightarrow v',J')$ and $A(v,J\\rightarrow v',J')$ are the observed line flux and the Einstein-$A$ radiative transition probability of the transition from level $(v,J)$ to $(v',J')$, respectively. The Einstein coefficients are adopted from Turner et al. (\\cite{tur77}) and Wolniewicz et al. (\\cite{wol98}). The transition energies we computed from level energies kindly provided by E. Roueff (1992, private communication).  A convenient way to visualize the level column densities is to divide them by the level degeneracy $g_J$, and plot this against the upper level energy $E_{\\rm u}(v,J)/k$; the degeneracy $g_J \\equiv g_s (2 J + 1)$, where $g_s = 3$ for ortho (odd $J$) H$_2$ and $g_s = 1$ for para (even $J$) H$_2$.  For the lines we observed toward Peak 1 we found that in such a ``Boltzmann diagram'' Fig.~\\ref{fig:boltz1} the level columns show a smooth distribution, where the level columns line up irrespective of their quantum numbers.  There is no sign of fluorescent excitation or of a deviation from the ortho-to-para H$_2$ ratio of three.   \\begin{figure}[htb] \\begin{center} \\includegraphics[width=1.0\\columnwidth]{figure5.ps} \\caption{Excitation diagram of the H$_2$ level column density distribution toward Peak 1, plotting the observed level columns (not corrected for extinction, divided by the level degeneracy) against level energy.} \\label{fig:boltz1} \\end{center} \\end{figure}  Fluorescently excited $\\rm H_2$, as seen in photodissociation regions (Timmermann et al. \\cite{tim96}) would produce a level distribution in which the ``rotational temperature'' derived from levels at given $v$ is lower than the ``vibrational temperature'' derived from levels of the same $J$ (e.g. Draine \\& Bertoldi \\cite{dra96}). Fluorescent excitation therefore shows a characteristic jigsaw distribution of the $v>1$ levels, unlike the smooth line-up we observed here, where $N/g$ appears not to depend on the state quantum number. Furthermore, fluorescently excited gas usually shows ortho-to-para ratios in vibrationally excited levels smaller than the total ortho-to-para ratio of the gas along the line of sight.  This is due to the enhanced self-shielding, and therefore reduced excitation rate, of the more abundant ortho-H$_2$ (Sternberg \\& Neufeld \\cite{ste99}).  \\subsection{Extinction} \\label{ex}  The shocks emitting the strong infrared lines in the OMC-1 outflow are deeply embedded in the molecular cloud, so that the emerging radiation suffers significant extinction. To correct the column densities derived from the H$_2$ emission line intensities, $N_{\\rm obs}(v,J)$, for this extinction, we need to know the proper extinction correction as a function of wavelength.  However, this interstellar infrared extinction curve is not well determined. Especially the depth and width of the silicate absorption features, centered at 9.7~$\\rm \\mu$m and 18~${\\rm \\mu m}$ are uncertain, and they could vary from region to region (Draine \\cite{dra89}). A further complication arises from the mixing of the emitting and absorbing gas, which we might expect in the outflow regions considering the complex spatial variation of the near-IR emission mapped in OMC-1 (Fig.~\\ref{schultz}), or in similar outflows such as DR21 and Cep A (Davis \\& Smith \\cite{dav96}; Goetz et al. \\cite{goe98}).  With enough redundancy in the information provided by the molecular and atomic lines in Peak~1, we are in principle able to estimate the average extinction along our line of sight as a function of wavelength.  However, the H~{\\sc i} recombination lines and the H$_2$ emission lines may not be tracing the same regions, and might therefore be subject to differing extinction. We therefore treat them separately.  To correct the $\\rm H_2$ line fluxes for extinction, we tried to derive the extinction from the $\\rm H_2$ lines directly (Bertoldi et al. \\cite{ber99}).  \\begin{figure}[htb] \\begin{center} \\includegraphics[width=1.0\\columnwidth]{figure6.ps} \\caption{Near- and mid-infrared extinction (Eq.~2) found from the relative intensities of the H$_2$ lines observed toward Peak~1. The curve was constructed using four free parameters for which values were derived that minimize the dispersion of the H$_2$ column density distribution (Fig.~5) for levels with $E/k < 16\\, 000$~K} \\label{extinction} \\end{center} \\end{figure}  An inspection of the excitation diagram Fig.~\\ref{fig:boltz1} derived from the line intensities ({\\em uncorrected} for extinction) shows that the column densities follow a smooth distribution, with no dependence on vibrational quantum number, and no sign of an ortho-to-para column density ratio different from three.  Transitions from a given state to different lower states produce lines with different wavelengths which suffer extinction.  Deviations from the expected line ratios of lines from the same state therefore yield the difference in extinction between the corresponding wavelengths of the lines. More generally, we can use this to estimate the extinction as a function of wavelength for a large set of lines, by minimizing the dispersion in the excitation diagram around a least-squares fit to the level columns.  We thereby assume that the dispersion in the column densities is partly due to extinction.  We constructed an extinction curve (Fig.  \\ref{extinction}) with four free parameters: the absolute normalization for a $A_\\lambda\\propto\\lambda^{\\rm -1.7}$ power law extinction curve from $\\rm 2.4\\, \\mu m$ to $6\\rm \\, \\mu m$, the width and depth of the water ice absorption feature at 3.1$\\,\\mu$m, and the depth of the $9.7\\, \\mu$m silicate absorption feature.  We fixed the width of the 9.7 and $\\rm 18 \\, \\mu m$ silicate features from calculations by Draine \\& Lee (\\cite{dra84}). The depth of the 18$\\, \\mu$m feature was taken to be 0.44 times that of the 9.7$\\rm \\,\\mu m $ feature, based on an average of previous estimates (Draine \\& Lee \\cite{dra84}; Pegourie \\& Papoular \\cite{peg85}; Volk \\& Kwok \\cite{vol88}; Bertoldi et al. in prep.). Calibration uncertainties and a small contribution of foreground fluorescently excited H$_2$ may give rise to a dispersion in the derived extinction corrections that is not due to extinction. The derived  curve should therefore be considered as very approximate. The extinction curve minimizing the dispersion in the excitation curve is shown in Fig.~\\ref{extinction}.  Explicitly it can be written \\begin{eqnarray} A(\\lambda) & = &  A_{\\rm K} (\\lambda/2.12)^{-1.7} + 0.58~e^{-22(\\lambda-3.05)^2} \\nonumber \\\\ &   &  {} + (1.35-0.08 A_{\\rm K}) \\left\\{e^{-[c_1 \\log(\\lambda/9.66)]^2} \\right. \\\\ &   &  \\left. {} + 0.44 ~e^{-[c_2 \\log(\\lambda/19)]^2} \\right\\}~, \\nonumber \\end{eqnarray} where $A_{\\rm K}=(1.0\\pm 0.1)$~mag is the implied extinction at 2.12~\\mum, and $c_1=14.3$ for $\\lambda<9.7$, $c_1=9.8$ for  $\\lambda>9.7$, $c_2=7.5$ for  $\\lambda<19$, $c_2=4.8$ for  $\\lambda>19$, and $\\lambda$ is given here in \\mum. The depth of the extinction minimum at 6.5~$\\mu$m is very uncertain, since it is not constrained by the inconsistent corrections derived from the four lines between 5.8 and 7.3~$\\mu$m. There is an indication that the minimum is at least as deep as our simple curve shows, but a much more careful analysis of the line fluxes would be necessary to reach a firm conclusion.  Atomic hydrogen recombination lines could offer another means to trace the extinction as a function of wavelength. We detected seventeen H~{\\sc i} recombination lines ranging in wavelength from 2.6 to 19~$\\rm \\mu m$. Since the relative emissivities are known from theory (Storey \\& Hummer \\cite{sto95}) and depend only mildly on the gas temperature and density, a comparison of the observed line intensities divided by their respective case B emissivities yields a measure for the differential extinction between the respective lines' wavelengths.  In Fig.~\\ref{recom} we plot against wavelength the observed line intensities, divided by their emissivities and normalized to this ratio for the H~{\\sc i} 8--5 transition line. The data points scatter around unity, which means that there is little if any differential extinction over this wavelength range.  A comparison with the distribution of intensities expected for an extinction curve with the shape we found from the H$_2$ lines reveals none of the prominent extinction features. A reasonable explanation is that the total extinction to the H~{\\sc i} emission region is very low, $A_{\\rm K}<0.3$ mag.  Although the errors are too large to constrain the exact value of the extinction, it is obvious that the H recombination lines are much less attenuated than the H$_2$ lines.  This suggests that the bulk of the atomic hydrogen emission arises in the foreground H~{\\sc ii} region, whereas the $\\rm H_2$ emitting region is more deeply embedded in the molecular cloud.  This conclusion agrees with the assessment of Everett et al. (\\cite{eve95}), who obtained $A_{\\rm J}=(0.38 \\pm 0.09)$~mag for the extinction shown by H recombination lines, but found $A_{\\rm J}=(2.15 \\pm 0.26)$~mag from the H$_2$ lines. With a $\\lambda^{-1.7}$ extinction law this corresponds to $A_{\\rm K}=(0.15 \\pm 0.04)$~mag and $A_{\\rm K}=(0.9 \\pm 0.1)$~mag, in good agreement with our results.   \\begin{figure*}[htb] \\begin{center} \\includegraphics[width=1.5\\columnwidth]{figure7.ps} \\caption{An attempt to derive the differential extinction in the H recombination line emission:  The line intensities divided by their respective emissivities are normalized by this value for the H 8--5 line, and then plotted against line wavelength.  With no differential extinction to the H 8--5 line, all data points should line up at unity value.  The broken lines show the expected distribution of values adopting the shape of the extinction curve derived from H$_2$ lines (Fig.\\ref{extinction}), with two different values of the absolute extinction at K.  The apparent distribution is consistent with an extinction of zero.} \\label{recom} \\end{center} \\end{figure*}   \\subsection{Fine structure lines} \\label{finestructure}  The observed atomic and ionic fine structure lines are valuable diagnostics. If highly-ionized species are found toward the shocked region -- which is well shielded from the ionizing radiation of the Trapezium stars -- they would indicate the presence of fast, ionizing J-shocks.  It is therefore interesting to disentangle the respective contributions to the fine structure line emission of the foreground H~{\\sc ii} region/PDR and the shocked gas of the OMC-1 outflow.  Table~\\ref{fine} lists the observed intensities of a number of fine structure lines we searched for toward Peak 1.  \\begin{table*} \\caption[]{ ISO-SWS Observations of Fine Structure Lines toward Orion Peak~1 } \\begin{tabular}{lrrrrccrlc} \\hline\\noalign{\\smallskip} \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{$\\lambda$} & \\multicolumn{1}{c}{$\\rm IP_l~^a$} & \\multicolumn{1}{c}{$\\rm IP_u~^b$} & \\multicolumn{1}{c}{$n_{\\rm crit}~^{\\rm c}$} & \\multicolumn{1}{c}{SWS} & \\multicolumn{1}{c}{$I_{\\rm obs}$} & \\multicolumn{1}{c}{S/N~$^{\\rm d}$ } & \\multicolumn{1}{c}{$A_\\lambda~^{\\rm e}$} & \\multicolumn{1}{c}{$I_{\\rm model}~^{\\rm f}$} \\\\ \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{[$\\rm \\mu$m]}    & \\multicolumn{1}{c}{[eV]}        & \\multicolumn{1}{c}{[eV]} & \\multicolumn{1}{c}{[cm$^{\\rm -3}$]} & \\multicolumn{1}{c}{AOT} & \\multicolumn{1}{c}{[erg~s$^{\\rm -1}$cm$^{\\rm -2}$sr$^{\\rm -1}$]} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{[mag]} & \\multicolumn{1}{c}{[erg~s$^{\\rm -1}$cm$^{\\rm -2}$sr$^{\\rm -1}$]} \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} [Ne~{\\sc iii}] & 36.009 & 40.96 & 63.45 & $5.52 \\times 10^4$ & 01 & $2.26 \\times 10^{-3}$ & 4.3  & 0.02 & $1.78 \\times 10^{-3}$ \\\\ {[Fe~{\\sc ii}]}  & 35.777 & 7.90  & 16.19 &                    & 01 & $1.22 \\times 10^{-3}$ & 2.0  & 0.02 & \\\\ {[Si~{\\sc ii}]}  & 34.814 & 8.15  & 16.35 & $3.67 \\times 10^{\\rm 5~g}$ & 01 & $1.43 \\times 10^{-2}$ & 19.1 & 0.03 & $3.16 \\times 10^{-4}$ \\\\ &        &       &       &                         & 02 & $1.59 \\times 10^{-2}$ & 58.8 &        & \\\\ {[S~{\\sc iii}]}  & 33.480 & 22.34 & 34.79 & $6.34 \\times 10^3$ & 01 & $2.16 \\times 10^{-2}$ & 46.0 & 0.03 & $1.20 \\times 10^{\\rm -2~h}$\\\\ {[Fe~{\\sc ii}]}  & 25.988 & 7.90  & 16.19 & $3.59 \\times 10^{\\rm 4~i}$ & 01 & $3.70 \\times 10^{-3}$ & 6.8  & 0.05 & \\\\ &        &       &       &                           & 02 & $1.76 \\times 10^{-3}$ &10.6  &      &  \\\\ {[S~{\\sc i}]}    & 25.249 & 0.0   & 10.36 & $4.2 \\times 10^{\\rm 4~j}$  & 01 &$1.19 \\times 10^{-2}$ & 19.5 & 0.35$~^{\\rm k}$ & \\\\ &        &       &       &                           & 02 & $1.17 \\times 10^{-2}$ & 20.5 &   &  \\\\ {[Fe~{\\sc iii}]} & 22.925 & 16.19 & 30.65 & $1.12 \\times 10^5$ & 01 & $4.15 \\times 10^{-4}$ & 2.1  & 0.07 & \\\\ {[S~{\\sc iii}]}  & 18.713 & 23.34 & 34.79 & $2.06 \\times 10^4$ & 01 & $3.61 \\times 10^{-2}$ & 307.0 & 0.09 & $3.63 \\times 10^{-2}$\\\\ {[Fe~{\\sc ii}]}  & 17.936 & 7.90  & 16.19 &                    & 01 & $2.56 \\times 10^{-4}$ & 2.2  & 0.09 & \\\\ {[P~{\\sc iii}]}  & 17.885 & 19.77 & 30.20 & $3.91 \\times 10^4$ & 01 & $3.09 \\times 10^{-4}$ & 2.9  & 0.09 & \\\\ {[Ne~{\\sc iii}]} & 15.555 & 40.96 & 63.45 & $2.70 \\times 10^5$ & 01 & $3.20 \\times 10^{-2}$ & 255.0 & 0.07 & $2.80 \\times 10^{-2}$\\\\ {[Ne~{\\sc ii}]}  & 12.814 & 21.56 & 40.96 & $6.54 \\times 10^5$ & 01 & $3.03 \\times 10^{-2}$ & 109.0 & 0.07 & $3.16 \\times 10^{-2}$ \\\\ &        &       &       &                    & 02 & $3.11 \\times 10^{-2}$ & 63.3 &       &  \\\\ {[S~{\\sc iv}]}$^{\\rm l}$ & 10.511 & 34.79 & 47.22 & $5.39 \\times 10^4$ & 01 & $9.69 \\times 10^{-3}$ & 550.0 & 0.19 & $2.40 \\times 10^{-2}$ \\\\ {[Ar~{\\sc iii}]} & 8.991  & 27.63 & 40.74 & $3.18 \\times 10^5$ & 01 & $9.29 \\times 10^{-3}$ & 323.0 & 0.18 & $1.38 \\times 10^{-2}$\\\\ {[Ar~{\\sc ii}]}  & 6.985  & 15.76 & 27.63 & $4.17 \\times 10^5$ & 01 & $2.92 \\times 10^{-3}$ & 148.0 & 0.02 & $2.82 \\times 10^{-3}$\\\\ {[Ni~{\\sc ii}]}$^{\\rm m}$ & 6.636  & 7.64  & 18.17 & $1.73 \\times 10^7$ & 01 & $1.26 \\times 10^{-3}$ &  47.8 & 0.02 & \\\\ \\noalign{\\smallskip}        \\hline \\noalign{\\smallskip} \\end{tabular}  $^{\\rm a}$ Lower ionization potential, IP$_{\\rm l}$, to produce the ion.\\\\ $^{\\rm b}$ Upper ionization potential, IP$_{\\rm u}$, of the next higher ionization stage.\\\\ $^{\\rm c}$ Collisions with electrons at $T = 10\\,000$~K, unless indicated.\\\\ $^{\\rm d}$ Calculated from the RMS noise within $\\rm \\sim 500 \\: km\\,s^{-1}$. \\\\ $^{\\rm e}$ From the extinction law of Bertoldi et al. (\\cite{ber99}) with $A_{\\rm K}=$0.15 mag.\\\\ $^{\\rm f}$ Best model by Rubin et al. (\\cite{rub91}) for a projected distance of $\\sim 86\\arcsec$ from $\\theta^1$~Ori~C. \\\\ $^{\\rm g}$ For collisions with electrons at $20\\,000$~K. Collisions with H atoms at 300 K yield $n_{\\rm crit} = 3.67 \\times 10^5 \\rm\\: cm^{-3}$.\\\\ $^{\\rm h}$ From the predicted [S~{\\sc iii}]18.7$\\mu$m line intensity and the [S~{\\sc iii}]18.7$\\mu$m/[S~{\\sc iii}]33.5$\\mu$m line ratio. \\\\ $^{\\rm i}$ Collisions with H atoms yield $n_{\\rm crit} = 2.24 \\times 10^6$ cm$^{-3}$.\\\\ $^{\\rm j}$ For collisions with $\\rm H^+$ ions. Collisions with H atoms at 300 K yield $n_{\\rm crit} = 1.5 \\times 10^6 \\rm\\: cm^{-3}$.\\\\ $^{\\rm k}$ Because [S~{\\sc i}]25.249$\\rm \\mu m$ probably arises from the more deeply embedded, shocked region, we adopt  $A_{\\rm K} = 1.0$ as for the H$_2$ emission, with the extinction curve of Fig.~\\ref{extinction}.\\\\ $^{\\rm l}$ Merged with H~12-8, but we estimate latter to contribute only $4\\times 10^{-5} \\: {\\rm erg~s^{-1}cm^{-2}sr^{-1}}$, adopting case-B emissivities. \\\\ $^{\\rm m}$ Merged with an $\\rm H_2O$ line at 6.6354 $\\rm \\mu m$.\\\\  \\label{fine} \\end{table*} %    A predominantly ionized medium is traced by species with ionization potentials larger than 13.6 eV.  From such ions a number of lines, [Ar~{\\sc ii}]6.9$\\mu$m, [Ar~{\\sc ii}]8.99$\\mu$m, [Ne~{\\sc ii}]12.8$\\mu$m, [Ne~{\\sc iii}]15.5$\\mu$m, [Ne~{\\sc iii}]36$\\mu$m, [S~{\\sc iii}]18.7$\\mu$m, and [S~{\\sc iv}]10.5$\\mu$m, are found in our ISO-SWS spectra, and their intensities can be compared with H~{\\sc ii} region models such as those computed by Rubin et al. (\\cite{rub91}). A comparison of the line intensities and their ratios to a blister H~{\\sc ii} region model with a star of $T_{\\rm eff}=37\\,000$~K and $\\log g = 4.0$, shows good agreement of all line intensities, except for that the models overestimate the [S~{\\sc iii}]18.7$\\mu$m/[S~{\\sc iii}]33.5$\\mu$m ratio by a factor 1.8.  The good agreement indicates that these lines may be predominantly produced in the foreground H~{\\sc ii} region, although a shock contribution of up to 30\\% cannot be excluded.  We can compare the Peak 1 fine structure line emission also with that seen toward the Orion Bar photodissociation region and ionization front, which is also irradiated by the Trapezium stars.  We know that here no fast shocks should contribute to the emission, and that the emission should be similar to that coming from the PDR in front of the OMC-1 outflow (Herrmann et al. \\cite{her97}).  In Table~\\ref{bar} we list line intensities we observed toward two positions on the Orion Bar: toward the ionization front at $5^{\\rm h} 35^{\\rm m} 19.31^{\\rm s}$, $-5 \\degr 24 \\arcmin 59.9\\arcsec$ (J2000), and toward the peak of the H$_2$ 1-0 S(1) emission at $5^{\\rm h} 35^{\\rm m} 20.31^{\\rm s}$, $-5 \\degr 25 \\arcmin 19.9\\arcsec$.  A comparison with the Peak 1 intensities shows that the intensities of most lines agree within a factor of a few, suggesting that ionic emission indeed arises in the foreground H~{\\sc ii} region.  The [P~{\\sc iii}]17.9$\\mu$m, [Fe~{\\sc iii}]22.9$\\mu$m, [Fe~{\\sc ii}]26$\\mu$m, and the [S~{\\sc iii}]33.5$\\mu$m lines were not included in the Rubin models, but their intensities are very similar toward the outflow and the Bar.  It is unclear where the [Fe~{\\sc ii}]26$\\mu$m emission comes from, though.  It could be produced either in the PDR or in the ionization front.  A detailed analysis of the [Fe~{\\sc ii}] emission will be subject of a subsequent publication (Bertoldi et al., in prep.).  \\begin{table*} \\caption[]{ Fine Structure Lines toward two Positions on the Orion Bar PDR } \\begin{tabular}{lrccrcrccrc} \\hline\\noalign{\\smallskip} \\hline\\noalign{\\smallskip} \\multicolumn{2}{c}{ } & \\multicolumn{4}{c}{\\rule[-2mm]{0mm}{6mm}Bar H$_2$ S(1)} & \\multicolumn{1}{c}{ } & \\multicolumn{4}{c}{Bar Br$\\gamma$} \\\\ \\cline{3-6} \\cline{8-11} \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{$\\lambda$} & \\multicolumn{1}{c}{SWS} & \\multicolumn{1}{c}{$I_{\\rm obs}$} & \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{SWS} & \\multicolumn{1}{c}{$I_{\\rm obs}$} & \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{ } \\\\ \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{[$\\rm \\mu$m]}    & \\multicolumn{1}{c}{AOT} & \\multicolumn{1}{c}{[erg~s$^{\\rm -1}$cm$^{\\rm -2}$sr$^{\\rm -1}$]} & \\multicolumn{1}{c}{\\raisebox{1.5ex}[-1.5ex]{S/N$~^{\\rm a}$}} & \\multicolumn{1}{c}{\\raisebox{1.5ex}[-1.5ex]{Bar/Pk1 $~^{\\rm b}$}} & \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{AOT} & \\multicolumn{1}{c}{[erg~s$^{\\rm -1}$cm$^{\\rm -2}$sr$^{\\rm -1}$]} & \\multicolumn{1}{c}{\\raisebox{1.5ex}[-1.5ex]{S/N}} & \\multicolumn{1}{c}{\\raisebox{1.5ex}[-1.5ex]{Bar/Pk1 $~^{\\rm b}$}} \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} {[Ne~{\\sc iii}]} & 36.009 & 01 & $1.75 \\times 10^{-3}$ &  12.9 &  0.8  & & 01 & $1.31 \\times 10^{-3}$ & 1.8  & 0.6 \\\\ {[Fe~{\\sc ii}]}  & 35.777 & 01 & $2.05 \\times 10^{-4}$ &   2.8 &  0.2  & &        &                       &      &     \\\\ {[Si~{\\sc ii}]}  & 34.814 & 01 & $1.38 \\times 10^{-2}$ & 135.0 &  1.1  & & 01 & $1.16 \\times 10^{-2}$ & 7.1  & 1.0 \\\\ &        &        &                       &       &       & & 02 & $1.13 \\times 10^{-2}$ & 63.7 & 0.7 \\\\ {[S~{\\sc iii}]}  & 33.480 & 01 & $2.70 \\times 10^{-2}$ &  88.4 &  1.3  & & 01 & $4.78 \\times 10^{-2}$ & 36.4 & 2.2 \\\\ {[Fe~{\\sc ii}]}  & 25.988 & 01 & $1.38 \\times 10^{-3}$ &  19.4 &  0.4  & & 01 & $9.58 \\times 10^{-4}$ &  4.2 & 0.3 \\\\ &      &        &                       &       &       & & 02 & $1.37 \\times 10^{-3}$ & 17.8 & 0.8 \\\\ {[S~{\\sc i}]}$^{\\rm d}$  & 25.249 & 01 & $ < 1.9 \\times 10^{-4}$ &   &$ < 0.02$ & & 01 & $< 2.1 \\times 10^{-4}$ &  &$< 0.02$  \\\\ {[Fe~{\\sc iii}]} & 22.925 & 01 & $4.82 \\times 10^{-4}$ &  10.6 &  1.2  & & 01 & $1.64 \\times 10^{-3}$ &  4.4 & 4.0 \\\\ &        &        &                       &       &       & & 02 & $9.60 \\times 10^{-4}$ & 15.6 & 2.3 \\\\ {[Ar~{\\sc iii}]} & 21.842 & 01 & $3.59 \\times 10^{-4}$ &   6.0 &       & &        &                       &      &     \\\\ {[S~{\\sc iii}]}  & 18.713 & 01 & $2.08 \\times 10^{-2}$ & 391.0 &  0.6  & & 01 & $5.19 \\times 10^{-2}$  & 249.0 & 1.4 \\\\ {[Fe~{\\sc ii}]}  & 17.936 & 01 & $4.22 \\times 10^{-5}$ &   2.1 &  0.1  & &        &                       &      &     \\\\ {[P~{\\sc iii}]}  & 17.885 & 01 & $1.56 \\times 10^{-4}$ &   6.6 &  0.5  & &        &                       &      &     \\\\ {[Ne~{\\sc iii}]} & 15.555 & 01 & $7.54 \\times 10^{-3}$ & 186.0 &  0.2  & & 01 & $2.62 \\times 10^{-2}$ & 575.0 & 0.8 \\\\ {[Ne~{\\sc ii}]}  & 12.814 & 01 & $1.49 \\times 10^{-2}$ &  76.5 &  0.5  & & 01 & $3.44 \\times 10^{-2}$ & 70.6 & 1.1 \\\\ {[S~{\\sc iv}]}$^{\\rm c}$ & 10.511 & 01 & $2.29 \\times 10^{-3}$ & 100.0 & 0.2   & & 01 & $8.38 \\times 10^{-3}$ & 55.0 & 0.9 \\\\ {[Ar~{\\sc iii}]} &  8.991 & 01 & $3.69 \\times 10^{-3}$ & 116.0 &  0.4  & & 01 & $1.13 \\times 10^{-2}$ & 194.0 & 1.2 \\\\ {[Ar~{\\sc ii}]}  & 6.985  & 01 & $1.90 \\times 10^{-3}$ &  18.7 &  0.7  & & 01 & $1.15 \\times 10^{-2}$ &  40.2 & 3.9 \\\\ {[Ni~{\\sc ii}]}$^{\\rm d}$  & 6.636  & 01 & $< 7.7 \\times 10^{-5}$ & & $< 0.06$ & & 01 & $< 1.4 \\times 10^{-4}$ &   & $< 0.1$ \\\\ \\noalign{\\smallskip}        \\hline \\noalign{\\smallskip} \\end{tabular}  $^{\\rm a}$ From the RMS noise within $\\rm \\sim 500 \\: km\\,s^{-1}$. \\\\ $^{\\rm b}$ Ratio of the respective intensities toward Bar H$_2$ S(1) and Peak 1, and Bar Br$\\gamma$ and Peak 1. If lines in different observing modes are available the ratio was calculated from  lines of the same mode.\\\\ $^{\\rm c}$ The [S~{\\sc iv}] line is merged with the H~12-8 line.\\\\ $^{\\rm d}$ The upper limit for the intensities is calculated from the $3 \\sigma$ flux density noise level at the respective wavelength times the width of one resolution element.  \\label{bar} \\end{table*} %   \\paragraph{Silicon:} Haas et al. (\\cite{haa91}) observed [Si~{\\sc ii}]34.8$\\mu$m strip maps across the OMC-1 outflow.  From the apparent peak of emission near IRc2 they concluded that about half of this emission must be due to the production and excitation of gas phase silicon in shocks.  In their preliminary reduction of a $\\sim 6 \\arcmin$ square map of [Si~{\\sc ii}], Stacey et al. (\\cite{sta95}) also find that the emission peaks toward the OMC-1 outflow, and this excess is consistent with that observed by Haas et al. (\\cite{haa91}).  Haas et al. find a surface flux density of $6\\times 10^{-3}\\rm ~erg~cm^{-2}s^{-1}sr^{-1}$ toward Peak 1, of which they attribute about half to an extended component, which most likely arises from the PDR lining the foreground H~{\\sc ii} region.  Haas et al. find that the flux is similar toward Peak~1 and the Orion Bar, which may well be due to limb brightening by a factor two of the PDR component at the Bar.  Our observations of [Si~{\\sc ii}]34.8$\\mu$m toward the Bar and Peak 1 yield fluxes twice as high as those of Haas.  This could be due either to a calibration error, or to beam dilution in the Haas et al. measurements. Either way, it seems that both the PDR and the shocks give rise to strong silicon emission, which for the PDR at least requires Si gas phase abundances of order 10\\% solar: the PDR models by Tielens \\& Hollenbach (\\cite{tie85b}) adopt a 2.2\\% solar gas phase Si abundance and predict about a quarter of the flux we could attribute to the PDR.  The silicon abundance can be enhanced in shocks by sputtering (Martin-Pintado et al. \\cite{mar92}; Caselli et al. \\cite{cas97}; Bachiller \\& Perez-Gutierrez \\cite{bac97}). Large gas phase silicon abundances are also found in other PDRs such as NGC 7023 (Fuente et al. \\cite{fue99}). The mechanism by which the abundance is enhanced in PDRs is still unclear, although photodesorption has been suggested (Walmsley et al. \\cite{wal99}). Strong silicate emission is however not a universal feature of PDRs: based on ISO observations of [Si~{\\sc ii}]34.8$\\mu$m toward NGC 2023 and a comparison with model calculations, Draine \\& Bertoldi (2000) report Si to be quite highly depleted in the NGC 2023 PDR.  \\paragraph{Other lines:} Of the other fine structure lines seen toward Peak 1, [Ni~{\\sc ii}]6.6$\\mu$m and [S~{\\sc i}]25$\\mu$m are not detected toward the Orion Bar. The [Ni~{\\sc ii}]6.6$\\mu$m line is confused with a water line, making it difficult to detect.  Although the [Fe~{\\sc ii}]18$\\mu$m and [Fe~{\\sc ii}]36$\\mu$m lines are marginally detected in both objects, they both appear to be an order of magnitude fainter toward the Bar than toward Peak~1. This suggests that shocks are more efficient in producing and exciting gas phase iron.  \\paragraph{Sulfur:} The strong [S~{\\sc i}]25$\\mu$m line emission is probably shock-excited. Burton et al. (\\cite{bur90a}) computed the [S~{\\sc i}] intensity in their PDR model for densities of $10^3$ to $\\rm 10^5 \\: cm^{-3}$ and radiation fields of $10^3$ to $10^5$ times the ambient interstellar field as $\\rm \\leq 10^{-5} \\: erg \\, s^{-1} \\, cm^{-2} \\, sr^{-1}$, which is three orders of magnitude below our observed [S~{\\sc i}] intensity. Both a J- or C-type shock could account for the [S~{\\sc i}] emission. But only a J-shock is able to produce both the [S~{\\sc i}] and the [Si~{\\sc ii}] line emission.  We compared the estimated shock contribution to the observed [Si~{\\sc ii}]34.8$\\mu$m flux of $\\rm \\sim 7 \\times 10^{-3} $ $\\rm erg \\, s^{-1} \\, cm^{-2} \\, sr^{-1}$ (Haas et al. \\cite{haa91}) and the [S~{\\sc i}]25$\\mu$m line flux to the J-shock model of Hollenbach \\& McKee (\\cite{hol89}). Both the relative and absolute [S~{\\sc i}] and [Si~{\\sc ii}] fluxes could be explained by shocks of high velocities, $v_{\\rm s} = {\\rm (85 \\pm 10) \\: km\\,s^{-1}}$, a pre-shock hydrogen nuclei density $ n_{\\rm H} = {\\rm (10^5 - 10^6) \\: cm^{-3}}$, and a beam filling factor $\\phi \\sim 3-4$. A beam-filling planar shock results in $\\phi = 1$, and a beam-filling spherical shock in $\\phi = 4$.  A shock contribution of 10 to 30\\% to the [Ne~{\\sc ii}]12.8$\\mu$m flux would also explain the observed [Ne~{\\sc ii}]/[Si~{\\sc ii}] and [Ne~{\\sc ii}]/[S~{\\sc i}] flux ratios.    \\subsection{Molecular hydrogen} \\label{molh2}  In the spectra shown in Figs.~\\ref{spectrum}, \\ref{spec}, and \\ref{multi}, we detected 56 different $\\rm H_2$ lines of pure rotational and rotation-vibrational transitions (Table~\\ref{h2_table}).  Pure rotational lines were detected ranging from the 0-0~S(1) to 0-0~S(25) transitions, which correspond to upper level energies $E(v,J)/k$ ranging from 1015~K to $42\\,500$~K. Adding a large number of vibration-rotational transition lines, we are able to study the excitation of the gas within the ISO aperture over an unprecedented range.  The H$_2$ 0-0 S(0) transition line was not detected from our observations with the medium resolution grating modes (SWS~01 and SWS~02, $R \\sim 1000-2000$).  Unfortunately, our observation with the Fabry-Perot did not cover a spectral range wide enough to detect a line with the expected width of $\\sim 60\\rm ~km ~s^{-1}$ (Nadeau \\& Geballe \\cite{nad79}; Brand et al. \\cite{bra89b}, Moorhouse et al. \\cite{moo90}; Chrysostomou et al. \\cite{chr97}).  However, the FP spectrum, shown in Fig.~\\ref{multi}, shows a line-like feature with a a narrow width of 12~km~s$^{-1}$, comparable to the spectral resolution in this observing mode.  This feature could be emission arising in the foreground photodissociation region bounding the Orion Nebula and the dense molecular cloud embedding the outflow.  \\begin{table*}[h] \\caption[]{ Summary of the ISO-SWS $\\rm H_2$ line observations at Orion Peak~1 } \\begin{tabular}{lrrlrcrrrr} \\hline\\noalign{\\smallskip} \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{$\\lambda$} & \\multicolumn{1}{c}{$E_{\\rm u}/k~^{\\rm a}$} & \\multicolumn{1}{c}{$A~^{\\rm b}$} & \\multicolumn{1}{c}{SWS} & \\multicolumn{1}{c}{$I_{\\rm obs}~^{\\rm c}$} & \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{$A_\\lambda~^{\\rm e}$} & \\multicolumn{1}{c}{$N_{\\rm u}~^{\\rm f}$} \\\\ \\multicolumn{1}{c}{\\raisebox{1.ex}[-1.ex]{line}} & \\multicolumn{1}{c}{[$\\rm \\mu m$]} & \\multicolumn{1}{c}{[K]} & \\multicolumn{1}{c}{[s$^{-1}$]} & \\multicolumn{1}{c}{AOT} & \\multicolumn{1}{c}{[erg~s$^{\\rm -1}$cm$^{\\rm -2}$sr$^{\\rm -1}$]} & \\multicolumn{1}{c}{\\raisebox{1.ex}[-1.ex]{S/N$^{\\rm d}$}} & \\multicolumn{1}{c}{[mag]} & \\multicolumn{1}{c}{[cm$^{-2}$]} \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} 0-0 S(0)  & 28.2188 &   509.9   & $2.94 \\times 10^{-11}$ & 01 & $< 7.90 \\times 10^{-4}$ & $< 3$.  & 0.27 & $< 6.17 \\times 10^{21}$ \\\\ 0-0 S(1)  & 17.0346 &  1015.0   & $4.76 \\times 10^{-10}$ & 01 & $ 1.34 \\times 10^{-3}$ &  27.8     & 0.53 & $4.92 \\times 10^{20}$  \\\\ &         &           &                        & 02 & $ 1.71 \\times 10^{-3}$ &  51.4     & 0.53 & $6.28 \\times 10^{20}$  \\\\ 0-0 S(2)  & 12.2785 &  1682.0   & $2.76 \\times 10^{-9}$  & 01 & $ 1.44 \\times 10^{-3}$ &  32.8     & 0.57 & $6.84 \\times 10^{19}$  \\\\ &         &           &                        & 02 & $ 1.78 \\times 10^{-3}$ &  12.6     & 0.57 & $8.46 \\times 10^{19}$  \\\\ 0-0 S(3)  &  9.6649 &  2503.4   & $9.84 \\times 10^{-9}$  & 01 & $ 4.08 \\times 10^{-3}$ & 200.0     & 1.35 & $8.82 \\times 10^{19}$  \\\\ &         &           &                        & 02 & $ 4.72 \\times 10^{-3}$ & 251.0     & 1.35 & $1.02 \\times 10^{20}$  \\\\ 0-0 S(4)  &  8.0258 &  3474.6   & $2.64 \\times 10^{-8}$  & 01 & $ 4.43 \\times 10^{-3}$ & 133.0     & 0.45 & $1.28 \\times 10^{19}$  \\\\ &         &           &                        & 02 & $ 4.84 \\times 10^{-3}$ & 231.0     & 0.45 & $1.40 \\times 10^{19}$  \\\\ 1-1 S(5)  &  7.2807 & $10\\,340.3$   & $5.44 \\times 10^{-8}$  & 01 & $ 2.48 \\times 10^{-4}$ &   6.7     & 0.18 & $2.48 \\times 10^{17}$  \\\\ 0-0 S(5)  &  6.9091 &  4586.7   & $5.88 \\times 10^{-8}$  & 01 & $ 1.09 \\times 10^{-2}$ &  63.7     & 0.15 & $9.31 \\times 10^{18}$  \\\\ &         &           &                        & 02 & $ 1.15 \\times 10^{-2}$ &  64.2     & 0.15 & $9.83 \\times 10^{18}$  \\\\ 0-0 S(6)  &  6.1089 &  5829.8   & $1.14 \\times 10^{-7}$  & 01 & $ 3.37 \\times 10^{-3}$ &  44.2     & 0.17 & $1.33 \\times 10^{18}$  \\\\ &         &           &                        & 02 & $ 3.03 \\times 10^{-3}$ &  58.3     & 0.17 & $1.20 \\times 10^{18}$  \\\\ 1-1 S(7)  &  5.8111 & $12\\,816.4$   & $1.82 \\times 10^{-7}$  & 01 & $ 2.48 \\times 10^{-4}$ &   8.1     & 0.18 & $5.91 \\times 10^{16}$  \\\\ 0-0 S(7)  &  5.5115 &  7196.6   & $2.00 \\times 10^{-7}$  & 01 & $ 9.99 \\times 10^{-3}$ & 113.0     & 0.20 & $2.09 \\times 10^{18}$  \\\\ &         &           &                        & 02 & $ 8.33 \\times 10^{-3}$ & 141.0     & 0.20 & $1.74 \\times 10^{18}$  \\\\ 0-0 S(8)  &  5.0528 &  8677.1   & $3.24 \\times 10^{-7}$  & 01 & $ 2.28 \\times 10^{-3}$ &  24.8     & 0.23 & $2.78 \\times 10^{17}$  \\\\ 1-1 S(9)$^{\\rm g}$  &  4.9533 & $15\\,725.5$   & $4.38 \\times 10^{-7}$  & 01 & $ 4.65 \\times 10^{-4}$ &   4.5     & 0.24 & $4.14 \\times 10^{16}$  \\\\ &         &           &                        & 06 & $ 2.81 \\times 10^{-4}$ &   4.1     & 0.24 & $2.50 \\times 10^{16}$  \\\\ 0-0 S(9)$^{\\rm g}$  &  4.6947 & $10\\,261.2$   & $4.90 \\times 10^{-7}$  & 01 & $ 5.09 \\times 10^{-3}$ &  15.7     & 0.26 & $3.92 \\times 10^{17}$  \\\\ &         &           &                        & 02 & $ 4.85 \\times 10^{-3}$ &  26.9     & 0.26 & $3.73 \\times 10^{17}$  \\\\ 1-1 S(11) &  4.4171 & $18\\,977.1$   & $8.42 \\times 10^{-7}$  & 01 & $ 2.13 \\times 10^{-4}$ &   4.6     & 0.29 & $9.21 \\times 10^{15}$  \\\\ 0-0 S(10) &  4.4096 & $11\\,940.2$   & $7.03 \\times 10^{-7}$  & 01 & $ 1.26 \\times 10^{-3}$ &  18.2     & 0.29 & $6.52 \\times 10^{16}$  \\\\ 3-2 O(7)  &  4.3298 & $19\\,092.2$   & $9.77 \\times 10^{-8}$  & 01 & $< 6.54 \\times 10^{-5}$ &  $< 3$     & 0.30 & $< 2.41 \\times 10^{16}$  \\\\ 2-2 S(13) &  4.3137 & $27\\,264.3$   & $1.14 \\times 10^{-6}$  & 01 & $< 7.47 \\times 10^{-5}$ & $< 3$    & 0.30 & $< 2.35 \\times 10^{15}$  \\\\ 0-0 S(11) &  4.1810 & $13\\,702.7$   & $9.64 \\times 10^{-7}$  & 01 & $ 2.39 \\times 10^{-3}$ &  31.7     & 0.32 & $8.77 \\times 10^{16}$  \\\\ &         &           &                        & 02 & $ 2.19 \\times 10^{-3}$ & 113.0     & 0.32 & $8.03 \\times 10^{16}$  \\\\ 1-0 O(8)  &  4.1622 &  9285.7   & $7.38 \\times 10^{-8}$  & 01 & $ 1.91 \\times 10^{-4}$ &   3.7     & 0.32 & $9.13 \\times 10^{16}$  \\\\ &         &           &                        & 02 & $ 1.10 \\times 10^{-4}$ &   7.9     & 0.32 & $5.26 \\times 10^{16}$  \\\\ 1-1 S(13) &  4.0675 & $22\\,516.4$   & $1.38 \\times 10^{-6}$  & 01 & $ 2.39 \\times 10^{-4}$ &  13.7     & 0.33 & $6.17 \\times 10^{15}$  \\\\ &         &           &                        & 02 & $ 1.80 \\times 10^{-4}$ &   8.7     & 0.33 & $4.55 \\times 10^{15}$  \\\\ 0-0 S(12) &  3.9968 & $15\\,538.5$   & $1.27 \\times 10^{-6}$  & 01 & $ 6.50 \\times 10^{-4}$ &  41.5     & 0.34 & $1.78 \\times 10^{16}$  \\\\ 1-1 S(14) &  3.9414 & $24\\,372.4$   & $1.69 \\times 10^{-6}$  & 01 & $ 1.03 \\times 10^{-4}$ &   6.5     & 0.35 & $2.09 \\times 10^{15}$  \\\\ 0-0 S(13) &  3.8464 & $17\\,437.7$   & $1.62 \\times 10^{-6}$  & 01 & $ 1.51 \\times 10^{-3}$ &  78.2     & 0.36 & $3.17 \\times 10^{16}$  \\\\ &         &           &                        & 02 & $ 1.40 \\times 10^{-3}$ &  92.7     & 0.36 & $2.94 \\times 10^{16}$  \\\\ 1-1 S(15) &  3.8404 & $26\\,257.2$   & $2.00 \\times 10^{-6}$  & 01 & $ 1.09 \\times 10^{-4}$ &   7.2     & 0.36 & $1.85 \\times 10^{15}$  \\\\ &         &           &                        & 02 & $ 1.51 \\times 10^{-4}$ &  10.8     & 0.36 & $2.57 \\times 10^{15}$  \\\\ 1-0 O(7)  &  3.8075 &  8364.9   & $1.06 \\times 10^{-7}$  & 01 & $ 1.43 \\times 10^{-3}$ &  99.3     & 0.37 & $4.57 \\times 10^{17}$  \\\\ 1-1 S(16) &  3.7602 & $28\\,199.5$   & $2.32 \\times 10^{-6}$  & 01 & $< 1.51 \\times 10^{-5}$ &   $< 3$    & 0.38 & $< 2.19 \\times 10^{14}$  \\\\ 2-1 O(6)  &  3.7236 & $13\\,150.2$   & $2.28 \\times 10^{-7}$  & 01 & $< 3.72 \\times 10^{-4}$ &$< 15.3 $   & 0.38 & $< 5.47 \\times 10^{16}$  \\\\ 0-0 S(14) &  3.7245 & $19\\,408.7$   & $2.41 \\times 10^{-6}$  & 01 & $< 3.72 \\times 10^{-4}$ &$< 15.3 $   & 0.38 & $< 5.18 \\times 10^{15}$  \\\\ 1-1 S(17) &  3.6979 & $30\\,156.2$   & $2.64 \\times 10^{-6}$  & 01 & $ 6.21 \\times 10^{-5}$ &   3.8     & 0.39 & $7.87 \\times 10^{14}$  \\\\ 3-2 O(5)  &  3.6630 & $17\\,811.7$   & $3.52 \\times 10^{-7}$  & 01 & $< 3.73 \\times 10^{-5}$ & $< 3$    & 0.39 & $< 3.53 \\times 10^{15}$  \\\\ 0-0 S(15) &  3.6263 & $21\\,408.6$   & $2.41 \\times 10^{-6}$  & 01 & $ 6.15 \\times 10^{-4}$ &  29.9     & 0.40 & $8.48 \\times 10^{15}$  \\\\ &         &           &                        & 02 & $ 6.49 \\times 10^{-4}$ &  31.2     & 0.40 & $8.94 \\times 10^{15}$  \\\\ 0-0 S(16) &  3.5470 & $23\\,451.6$   & $2.83 \\times 10^{-6}$  & 01 & $ 2.81 \\times 10^{-4}$ &   7.2     & 0.42 & $3.28 \\times 10^{15}$  \\\\ 1-0 O(6)  &  3.5007 &  7583.7   & $1.50 \\times 10^{-7}$  & 01 & $ 8.37 \\times 10^{-4}$ &  54.5     & 0.43 & $1.84 \\times 10^{17}$  \\\\ 0-0 S(17) &  3.4856 & $25\\,537.8$   & $3.26 \\times 10^{-6}$  & 01 & $ 3.35 \\times 10^{-4}$ &  23.1     & 0.44 & $3.39 \\times 10^{15}$  \\\\ 2-1 O(5)  &  3.4384 & $12\\,550.2$   & $3.18 \\times 10^{-7}$  & 01 & $< 4.31 \\times 10^{-4}$ &$< 22.8$ & 0.46 & $< 4.51 \\times 10^{16}$  \\\\ 0-0 S(18) & 3.4384 & $27\\,638.4$   & $3.68 \\times 10^{-6}$  & 01 & $< 4.31 \\times 10^{-4}$ &$< 22.8$ & 0.46 & $< 3.89 \\times 10^{15}$  \\\\ 0-0 S(19) &  3.4039 & $29\\,767.7$   & $4.08 \\times 10^{-6}$  & 01 & $ 1.62 \\times 10^{-4}$ &  15.1     & 0.48 & $1.34 \\times 10^{15}$  \\\\ &         &           &                        & 02 & $ 1.13 \\times 10^{-4}$ &  17.2     & 0.48 & $9.31 \\times 10^{14}$  \\\\ \\noalign{\\smallskip}        \\hline \\noalign{\\smallskip} \\end{tabular} \\label{h2_table} \\end{table*} %   \\addtocounter{table}{-1} \\begin{table*}[h] \\caption[]{ -Continued } \\begin{tabular}{lrrlrcrrrr} \\hline\\noalign{\\smallskip} \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{$\\lambda$} & \\multicolumn{1}{c}{$E_{\\rm u}/k~^{\\rm a}$} & \\multicolumn{1}{c}{$A~^{\\rm b}$} & \\multicolumn{1}{c}{SWS} & \\multicolumn{1}{c}{$I_{\\rm obs}~^{\\rm c}$} & \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{$A_\\lambda~^{\\rm e}$} & \\multicolumn{1}{c}{$N_{\\rm u}~^{\\rm f}$} \\\\ \\multicolumn{1}{c}{\\raisebox{1.ex}[-1.ex]{line}} & \\multicolumn{1}{c}{[$\\rm \\mu m$]} & \\multicolumn{1}{c}{[K]} & \\multicolumn{1}{c}{[s$^{-1}$]} & \\multicolumn{1}{c}{AOT} & \\multicolumn{1}{c}{[erg~s$^{\\rm -1}$cm$^{\\rm -2}$sr$^{\\rm -1}$]} & \\multicolumn{1}{c}{\\raisebox{1.ex}[-1.ex]{S/N$~^{\\rm d}$}} & \\multicolumn{1}{c}{[mag]} & \\multicolumn{1}{c}{[cm$^{-2}$]} \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} 3-2 O(4)  &  3.3958 & $17\\,380.1$   & $4.87 \\times 10^{-7}$  & 01 & $ 4.88 \\times 10^{-4}$ &   6.5     & 0.49 & $3.38 \\times 10^{15}$  \\\\ 0-0 S(20) &  3.3809 & $31\\,898.5$   & $4.45 \\times 10^{-6}$  & 01 & $< 1.43 \\times 10^{-5}$ &  $< 3$     & 0.50 & $< 1.09 \\times 10^{14}$  \\\\ 0-0 S(21) &  3.3689 & $34\\,040.8$   & $4.78 \\times 10^{-6}$  & 01 & $ 2.47 \\times 10^{-5}$ &   4.0     & 0.52 & $1.77 \\times 10^{14}$  \\\\ 0-0 S(22) &  3.3663 & $36\\,149.7$   & $5.06 \\times 10^{-6}$  & 01 & $< 1.91 \\times 10^{-5}$ &$< 3$    & 0.52 & $< 1.30 \\times 10^{14}$  \\\\ 0-0 S(23) &  3.3718 & $38\\,299.5$   & $5.27 \\times 10^{-6}$  & 01 & $< 1.91 \\times 10^{-5}$ & $< 3$   & 0.51 & $< 1.24 \\times 10^{14}$  \\\\ 0-0 S(24) &  3.3876 & $40\\,419.6$   & $5.42 \\times 10^{-6}$  & 01 & $< 1.42 \\times 10^{-5}$ &  $< 3$    & 0.50 & $< 8.88 \\times 10^{13}$  \\\\ 0-0 S(25) &  3.4108 & $42\\,515.1$   & $5.50 \\times 10^{-6}$  & 01 & $< 3.33 \\times 10^{-5}$ & $< 3$  & 0.48 & $< 2.03 \\times 10^{14}$  \\\\ &         &           &                        & 02 & $ 2.42 \\times 10^{-5}$ &   4.1     & 0.48 & $1.48 \\times 10^{14}$  \\\\ 0-0 S(26) &  3.4417 & $44\\,573.2$   & $5.51 \\times 10^{-6}$  & 01 & $< 3.27 \\times 10^{-5}$ & $< 3$ & 0.46 & $< 1.97 \\times 10^{14}$  \\\\ 0-0 S(27) &  3.4855 & $46\\, 650.3$  & $5.43 \\times 10^{-6}$  & 01 & $< 3.27 \\times 10^{-5}$ & $< 3$ & 0.44 & $< 1.82 \\times 10^{14}$  \\\\ 0-0 S(28) &  3.5375 & $48\\,640.3$   & $5.28 \\times 10^{-6}$  & 01 & $< 3.86 \\times 10^{-5}$ & $< 3$   & 0.42 & $< 2.41 \\times 10^{14}$  \\\\ 0-0 S(29) &  3.5996 & $50\\,619.9$   & $5.04 \\times 10^{-6}$  & 01 & $< 4.56 \\times 10^{-5}$ & $< 3$  & 0.41 & $< 3.00 \\times 10^{14}$  \\\\ 1-0 O(5)  &  3.2350 & 6950.6   & $2.09 \\times 10^{-7}$  & 01 & $ 3.24 \\times 10^{-3}$ &  221.0    & 0.76 & $6.39 \\times 10^{17}$  \\\\ &         &           &                        & 02 & $ 3.05 \\times 10^{-3}$ &  218.0    & 0.76 & $6.02 \\times 10^{17}$  \\\\ &         &           &                        & 06 & $ 3.05 \\times 10^{-3}$ &  144.0    & 0.76 & $6.02 \\times 10^{17}$  \\\\ 2-1 O(4)  &  3.1899 & $12\\,094.1$   & $4.41 \\times 10^{-7}$  & 01 & $ 1.25 \\times 10^{-4}$ &   12.4    & 0.88 & $1.28 \\times 10^{16}$  \\\\ 3-2 O(3)  &  3.1637 & $17\\,092.3$   & $7.04 \\times 10^{-7}$  & 01 & $ 4.08 \\times 10^{-5}$ &    3.8    & 0.94 & $2.76 \\times 10^{15}$  \\\\ 1-0 O(4)  &  3.0039 &  6471.5   & $2.90 \\times 10^{-7}$  & 01 & $ 1.28 \\times 10^{-3}$ &  119.0    & 1.11 & $2.32 \\times 10^{17}$  \\\\ &         &           &                        & 02 & $ 1.20 \\times 10^{-3}$ &   70.5    & 1.11 & $2.18 \\times 10^{17}$  \\\\ &         &           &                        & 06 & $ 1.34 \\times 10^{-3}$ &   49.6    & 1.11 & $2.43 \\times 10^{17}$  \\\\ 2-1 O(3)  &  2.9741 & $11\\,789.1$   & $6.40 \\times 10^{-7}$  & 01 & $ 3.43 \\times 10^{-4}$ &   24.4    & 1.07 & $2.71 \\times 10^{16}$  \\\\ 3-2 O(2)  &  2.9620 & $16\\,948.5$   & $1.41 \\times 10^{-6}$  & 01 & $< 1.98 \\times 10^{-5}$ &  $< 3$    & 1.06 & $< 6.96 \\times 10^{14}$  \\\\ 2-1 Q(13) &  2.9061 & $23\\,926.4$   & $2.22 \\times 10^{-7}$  & 01 & $< 4.37 \\times 10^{-5}$ & $< 3$    & 0.95 & $< 8.70 \\times 10^{15}$  \\\\ 3-2 Q(7)  &  2.8250 & $20\\,861.9$   & $3.58 \\times 10^{-7}$  & 01 & $< 4.84 \\times 10^{-5}$ & $< 3$    & 0.80 & $< 5.07 \\times 10^{15}$  \\\\ 1-0 O(3)  &  2.8025 &  6149.2   & $4.23 \\times 10^{-7}$  & 01 & $ 6.17 \\times 10^{-3}$ &  388.0    & 0.77 & $5.27 \\times 10^{17}$  \\\\ &         &           &                        & 02 & $ 5.20 \\times 10^{-3}$ &  287.0    & 0.77 & $4.44 \\times 10^{17}$  \\\\ 2-1 O(2)  &  2.7862 & $11\\,635.2$   & $1.29 \\times 10^{-6}$  & 01 & $ 1.11 \\times 10^{-4}$ &    6.5    & 0.75 & $3.04 \\times 10^{15}$  \\\\ 3-2 Q(5)  &  2.7692 & $19\\,092.2$   & $3.98 \\times 10^{-7}$  & 01 & $ 9.81 \\times 10^{-5}$ &    5.9    & 0.74 & $8.52 \\times 10^{15}$  \\\\ 3-2 Q(3)  &  2.7312 & $17\\,811.7$   & $4.41 \\times 10^{-7}$  & 01 & $< 6.55 \\times 10^{-5}$ & $< 3$    & 0.71 & $< 4.94 \\times 10^{15}$  \\\\ 1-0 Q(13) &  2.7269 & $18\\,977.1$   & $1.61 \\times 10^{-7}$  & 01 & $ 1.10 \\times 10^{-4}$ &    3.3    & 0.71 & $2.27 \\times 10^{16}$  \\\\ 3-2 Q(2)  &  2.7186 & $17\\,394.5$   & $4.84 \\times 10^{-7}$  & 01 & $< 2.63 \\times 10^{-4}$ & $<   8.4 $ & 0.71 & $< 1.79 \\times 10^{16}$  \\\\ 2-1 Q(9)  &  2.7200 & $18\\,099.5$   & $3.03 \\times 10^{-7}$  & 01 & $< 2.63 \\times 10^{-4}$ & $<   8.4 $ & 0.71 & $< 2.87 \\times 10^{16}$  \\\\ 3-2 Q(1)  &  2.7102 & $17\\,092.3$   & $6.86 \\times 10^{-7}$  & 01 & $< 5.43 \\times 10^{-5}$ &  $< 3$    & 0.70 & $< 2.60 \\times 10^{15}$  \\\\ 2-1 Q(8)  &  2.6850 & $16\\,876.5$   & $3.22 \\times 10^{-7}$  & 01 & $< 5.29 \\times 10^{-5}$ & $< 3$  & 0.70 & $< 5.32 \\times 10^{15}$  \\\\ 2-1 Q(7)  &  2.6538 & $15\\,768.7$   & $3.40 \\times 10^{-7}$  & 01 & $ 3.32 \\times 10^{-4}$ &   10.3    & 0.70 & $3.13 \\times 10^{16}$  \\\\ 1-0 Q(11) &  2.6350 & $15\\,725.5$   & $1.87 \\times 10^{-7}$  & 01 & $ 3.10 \\times 10^{-4}$ &    9.8    & 0.70 & $5.29 \\times 10^{16}$  \\\\ &         &           &                        & 02 & $ 1.86 \\times 10^{-4}$ &    5.6    & 0.70 & $3.17 \\times 10^{16}$  \\\\ 2-1 Q(5)  &  2.6040 & $13\\,889.7$    & $3.74 \\times 10^{-7}$  & 01 & $ 6.54 \\times 10^{-4}$ &   17.2    & 0.71 & $5.55 \\times 10^{16}$  \\\\ &         &            &                        & 02 & $ 3.58 \\times 10^{-4}$ &   40.5    & 0.71 & $3.04 \\times 10^{16}$  \\\\ 1-0 Q(10) &  2.5954 & $14\\,220.6$    & $1.99 \\times 10^{-7}$  & 01 & $ 1.55 \\times 10^{-4}$ &    9.1    & 0.72 & $2.47 \\times 10^{16}$  \\\\ &         &            &                        & 02 & $ 1.59 \\times 10^{-4}$ &    9.3    & 0.72 & $2.53 \\times 10^{16}$  \\\\ 2-1 Q(4)  &  2.5850 & $13\\,150.2$    & $2.65 \\times 10^{-7}$  & 01 & $ 1.60 \\times 10^{-4}$ &    7.8    & 0.72 & $1.91 \\times 10^{16}$  \\\\ 2-1 Q(3)  &  2.5698 & $12\\,550.2$    & $4.12 \\times 10^{-7}$  & 01 & $ 4.42 \\times 10^{-4}$ &   13.8    & 0.72 & $3.40 \\times 10^{16}$  \\\\ 1-0 Q(9)  &  2.5600 & $12\\,816.4$    & $2.12 \\times 10^{-7}$  & 01 & $< 8.27 \\times 10^{-4}$ & $<  4.8 $   & 0.73 & $< 1.24 \\times 10^{17}$  \\\\ &         &            &                        & 02 & $< 7.85 \\times 10^{-4}$ & $< 59.5  $  & 0.73 & $< 1.17 \\times 10^{17}$  \\\\ 2-1 Q(2)  &  2.5585 & $12\\,094.1$    & $4.50 \\times 10^{-7}$  & 01 & $< 8.27  \\times 10^{-4}$ & $< 4.8 $    & 0.73 & $<5.82 \\times 10^{16}$  \\\\ &         &         &                           & 02 & $< 7.85 \\times 10^{-4}$ & $< 59.5 $  & 0.73 & $< 5.53 \\times 10^{16}$  \\\\ 2-1 Q(1)  &  2.5510 & $11\\,789.1$    & $6.37 \\times 10^{-7}$  & 01 & $ 4.62 \\times 10^{-4}$ &   19.2    & 0.73 & $2.30 \\times 10^{16}$  \\\\ &         &            &                        & 02 & $ 4.25 \\times 10^{-4}$ &   43.4    & 0.73 & $2.11 \\times 10^{16}$  \\\\ 4-3 S(1)  & 2.5415  & $22\\,761.0$    & $4.49 \\times 10^{-7}$  & 02 & $6.98 \\times 10^{-5}$  &   4.5     & 0.74 & $4.93 \\times 10^{15}$  \\\\ 1-0 Q(8)$^{\\rm h}$  &  2.5278 & $11\\,521.5$    & $2.23 \\times 10^{-7}$  & 01 & $ 3.28 \\times 10^{-4}$ &   10.1    & 0.74 & $4.66 \\times 10^{16}$  \\\\ &         &            &                        & 02 & $ 3.30 \\times 10^{-4}$ &   23.3    & 0.74 & $4.69 \\times 10^{16}$  \\\\ 1-0 Q(7)  &  2.5001 & $10\\,340.3$    & $2.34 \\times 10^{-7}$  & 01 & $ 1.76 \\times 10^{-3}$ &   95.2    & 0.76 & $2.38 \\times 10^{17}$  \\\\ 1-0 Q(6)  &  2.4755 &  9285.7    & $2.45 \\times 10^{-7}$  & 01 & $ 8.49 \\times 10^{-4}$ &   29.9    & 0.77 & $1.10 \\times 10^{17}$  \\\\ \\noalign{\\smallskip}        \\hline \\noalign{\\smallskip} \\end{tabular} \\label{h2_table} \\end{table*} %  \\addtocounter{table}{-1} \\begin{table*}[h] \\caption[]{ -Continued } \\begin{tabular}{lrrlrcrrrr} \\hline\\noalign{\\smallskip} \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{$\\lambda$} & \\multicolumn{1}{c}{$E_{\\rm u}/k~^{\\rm a}$} & \\multicolumn{1}{c}{$A~^{\\rm b}$} & \\multicolumn{1}{c}{SWS} & \\multicolumn{1}{c}{$I_{\\rm obs}~^{\\rm c}$} & \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{$A_\\lambda~^{\\rm e}$} & \\multicolumn{1}{c}{$N_{\\rm u}~^{\\rm f}$} \\\\ \\multicolumn{1}{c}{\\raisebox{1.ex}[-1.ex]{line}} & \\multicolumn{1}{c}{[$\\rm \\mu m$]} & \\multicolumn{1}{c}{[K]} & \\multicolumn{1}{c}{[s$^{-1}$]} & \\multicolumn{1}{c}{AOT} & \\multicolumn{1}{c}{[erg~s$^{\\rm -1}$cm$^{\\rm -2}$sr$^{\\rm -1}$]} & \\multicolumn{1}{c}{\\raisebox{1.ex}[-1.ex]{S/N$^{\\rm d}$}} & \\multicolumn{1}{c}{[mag]} & \\multicolumn{1}{c}{[cm$^{-2}$]} \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} 1-0 Q(5)  &  2.4548 &  8364.9    & $2.55 \\times 10^{-7}$  & 01 & $ 3.68 \\times 10^{-3}$ &  209.0    & 0.78 & $4.60 \\times 10^{17}$  \\\\ 1-0 Q(4)  &  2.4375 &  7583.7    & $2.65 \\times 10^{-7}$  & 01 & $ 1.61 \\times 10^{-3}$ &   76.9    & 0.79 & $1.94 \\times 10^{17}$  \\\\ 1-0 Q(3)  &  2.4237 &  6950.6    & $2.78 \\times 10^{-7}$  & 01 & $ 5.52 \\times 10^{-3}$ &  176.0    & 0.80 & $6.34 \\times 10^{17}$  \\\\ 1-0 Q(2)  &  2.4134 &  6471.5    & $3.03 \\times 10^{-7}$  & 01 & $ 1.91 \\times 10^{-3}$ &   56.3    & 0.80 & $2.01 \\times 10^{17}$  \\\\ 1-0 Q(1)  &  2.4066 &  6149.2    & $4.29 \\times 10^{-7}$  & 01 & $ 6.31 \\times 10^{-3}$ &  195.0    & 0.81 & $4.71 \\times 10^{17}$  \\\\ \\noalign{\\smallskip}        \\hline \\noalign{\\smallskip} \\end{tabular}  $^{\\rm a}$ The upper level energies were kindly provided by Roueff (1992, private communication). \\\\ $^{\\rm b}$ The Einstein coefficients are taken from  Turner et al. (\\cite{tur77}) and Wolniewicz et al. (\\cite{wol98}). \\\\ $^{\\rm c}$ The upper limit for the intensities is calculated from the $3 \\sigma$ flux density noise level at the respective wavelength, times the width of one resolution element. In a few cases of merged lines it was not possible to derive the individual line intensities. There the measured intensity of the combined structure -- although the S/N ratio is $> 3$ -- is also indicated as an upper limit for each of the components. \\\\ $^{\\rm d}$ Calculated from the RMS noise within $\\rm \\sim 500 \\: km\\,s^{-1}$. \\\\ $^{\\rm e}$ From the extinction curve shown in Fig.~\\ref{extinction}, derived from the H$_2$ lines. \\\\ $^{\\rm f}$ Extinction-corrected upper level column density. \\\\ $^{\\rm g}$ The 1-1 S(9) and 0-0 S(9) lines are merged with CO lines. To derive the H$_2$ line intensities, the CO lines and the background were fit by sine functions plus a first order polynomial, whereas the H$_2$ lines were fit by a Gaussian. \\\\ $^{\\rm h}$ The 1-0 Q(8) line is merged with H 16-5 line. The respective intensities were derived by fitting two gaussian functions to the combined structure. \\label{h2_table} \\end{table*} %   \\subsubsection{Contribution from the foreground PDR}  The line emission toward Peak 1 must include some contribution from the photodissociation region bordering the foreground Orion Nebula H~{\\sc ii} region.  Garden (\\cite{gar86}) produced a $\\rm H_2$~1-0~S(1) map which covers OMC-1, the Trapezium, and the Orion Bar PDR.  Following Burton \\& Puxley (\\cite{bur90b}), we estimate that the extended fluorescent $\\rm H_2$ flux should amount to about 5\\% of the total $\\rm H_2$ emission toward Peak~1 over the SWS aperture.  For an additional estimate of the expected PDR contribution we can compare the total H$_2$ luminosity toward Peak 1 to that toward the Orion Bar, a PDR observed nearly edge-on south-east of the Trapezium. We did observe the Bar with the ISO-SWS (Bertoldi et al., in prep.), and find that the total H$_2$ emission here amounts to 0.008~ $\\rm erg~s^{-1}cm^{-2}sr^{-1}$, compared to the 0.28~$\\rm erg~s^{-1}cm^{-2}sr^{-1}$ toward Peak 1.  Since the Bar is the brightest PDR emission peak in the Orion Nebula, we see that the $\\rm H_2$ emission from the PDR toward Peak 1 is probably small compared with the emission arising from the deeply embedded outflow.  \\subsubsection{Excitation of molecular hydrogen}  From the line intensities we derived observed column densities of the levels from which these transitions arise (Eq.\\ref{eq:col}).  We correct these values for extinction with the curve we derived in Sec.~\\ref{ex} (Fig.~\\ref{extinction}), to obtain the inherent level column densities \\begin{equation} N(v,J) = N_{\\rm obs}(v,J)~~ 10^{0.4A(\\lambda)}. \\end{equation} The resulting excitation (Boltzmann) diagram is shown in Fig.~\\ref{excit}.  \\begin{figure*}[htb] \\begin{center} \\includegraphics[width=2.\\columnwidth]{figure8.ps} \\caption{Extinction-corrected, observed H$_2$ level column densities, divided by their degeneracy, plotted against the upper level energy $E(v,J)$. Vibrational levels are distinguished by different symbols: squares, triangles, diamonds, $+$, and $\\times$ represent $v=0,$ 1, 2, 3, and 4, respectively. The dotted line represents the fit Eq.~\\ref{fit}. Error bars represent $1\\sigma$ flux uncertainties of the line flux integrations, but do not include $\\sim$30\\% calibration uncertainties. Some lines were measured in two different AOTs, with the corresponding column densities both shown.} \\label{excit} \\end{center} \\end{figure*}  The lack of signs of fluorescent excitation in the level columns suggest that the molecules might be mostly thermally excited.  An H$_2$ column $N_{\\rm H_2,tot}$ in statistical (thermodynamic) equilibrium at a single kinetic temperature $T$ would yield a level distribution \\begin{equation} \\label{single_boltz} { N(v,J) \\over g_J} ~=~  N_{\\rm H_2,tot} \\; { e^{- E(v,J)/k T}\\over \\sum_{v',J'} g_{J'} \\, e^{- E(v',J')/k T} }, \\end{equation} which produces a straight line in the Boltzmann (excitation) diagram Fig.~\\ref{excit}.  An excitation temperature function, $T_{\\rm ex}(E)$, can be assigned to the level distributions at each level energy, $E(v,J)$, by computing the inverse of the derivative of the line which best fits $ \\ln[N(v,J)/g_J]$ as a function of $E(v,J)/k$.  Near the lowest energy levels, $T_{\\rm ex} \\sim 600$~K, whereas at $E(v,J)/k \\geq 14\\,000$~K, the excitation temperature rises to $\\sim 3200$~K.  To describe the range of excitation temperatures, we decomposed the distribution of column densities to a sum of five Boltzmann distributions of different excitation temperatures: \\begin{equation} \\label{fit} N(v,J)/ g_J = \\sum_{i=1}^5  C_i \\; e^{- E(v,J)/kT_{{\\rm ex},i} }, \\end{equation} where we chose $T_{{\\rm ex},i} = (628, 800, 1200, 1800, 3226)$~K, and the $C_i$ (see Table \\ref{column}) were determined by a least-squares-fit to the observed level columns.  In Fig.~\\ref{excit} the dotted line shows the five-component fit.  From this fit we can also compute the total warm $\\rm H_2$ column density, by summing the column densities over {\\it all} levels following the interpolated level column distribution: \\begin{eqnarray} \\label{n_tot} N_{\\rm H_2,tot} &  =  & \\sum_{v,J} \\left[ \\frac{N(v,J)}{g_J} \\right] \\: g_J \\nonumber \\\\ &  =  & \\sum_{v,J} \\sum_{i=1}^5 g_J \\: C_i \\; e^{-E(v,J)/k T_{{\\rm ex},i}} \\nonumber \\\\ &  =  &  (1.9\\pm 0.5) \\times 10^{21} \\: {\\rm cm^{-2}} . \\end{eqnarray} Adopting a distance of 450~pc (Genzel \\& Stutzki \\cite{gen89}), this column corresponds to a warm H$_2$ mass of $(0.06\\pm 0.015)~{\\rm M}_{\\sun}$ within the ISO-SWS aperture.  By summing from $J = 0$, we extrapolated the observed H$_2$ $(v=0,J\\geq 3)$ level populations to the unobserved $(v=0,J = 0, 1, 2)$ levels. Note that thereby we estimate the total {\\it warm} $\\rm H_2$ column density, but we do not account for the {\\it total} $\\rm H_2$ column along the line of sight, which includes an additional $\\rm \\approx 10^{22} \\: cm^{-2}$ cold gas from the molecular cloud which embeds the outflow. Most of this cold $\\rm H_2 $ resides in the ground states $J = 0$ and $J = 1$, and does not contribute to the emission observed from the shock-excited gas in the outflow.  By changing the order of summation in Eq.~\\ref{n_tot} we can compute the column densities corresponding to the five excitation temperature components, $N_{{\\rm H_2,}i}$ (Table~\\ref{column}), such that \\begin{equation} \\label{t_components} N_{\\rm H_2,tot} = \\sum _{i=1}^5 \\sum_{v,J} g(J) \\: C_i \\; e^{-E(v,J)/k \\, T_{{\\rm ex,}i}} = \\sum _{i=1}^5 N_{{\\rm H_2,}i}. \\end{equation}  \\begin{table} % \\caption[]{ } \\begin{tabular}{rccr} \\hline\\noalign{\\smallskip} \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{\\rule[-2mm]{0mm}{5mm}$T_{{\\rm ex,}i}$} & \\multicolumn{1}{c}{$C_i$} & \\multicolumn{1}{c}{$N_{\\rm H_{\\rm 2,i}}$} & \\multicolumn{1}{r}{fraction of}\\\\ \\multicolumn{1}{r}{(K)} & \\multicolumn{1}{c}{(cm$^{\\rm -2}$)} & \\multicolumn{1}{c}{(cm$^{\\rm -2}$)} & \\multicolumn{1}{c}{warm $\\rm H_2$} \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} 628  & $8.80 \\times 10^{19}$ & $1.37 \\times 10^{21}$ & 72.2\\%  \\\\ 800  & $2.01 \\times 10^{19}$ & $3.96 \\times 10^{20}$ & 20.9\\%  \\\\ 1200  & $3.62 \\times 10^{18}$ & $1.07 \\times 10^{20}$ &  5.7\\%  \\\\ 1800  & $3.83 \\times 10^{17}$ & $1.76 \\times 10^{19}$ &  0.9\\%  \\\\ 3226  & $7.01 \\times 10^{16}$ & $6.82 \\times 10^{18}$ &  0.4\\%  \\\\ \\hline \\noalign{\\smallskip} \\end{tabular} \\label{column} \\end{table} %  Figure~\\ref{temp} shows the corresponding cummulative column density, $N_{\\rm H_2}(T_{ex}>T)$, plotted against $T$. \\begin{figure}[b] \\begin{center} \\includegraphics[width=1.0\\columnwidth]{figure9.ps} \\caption{ Cummulative column density of H$_2$ with excitation temperatures larger than a given value of $T$, observed within the ISO aperture. Derived from the five-component decomposition listed in Table 4.  } \\label{temp} \\end{center} \\end{figure} With the interpolated excitation distribution Eq.~\\ref{fit} the column densities of all $\\rm H_2$ energy levels can be estimated, even those from which no lines were observed. Then the total H$_2$ rovibrational emission from the electronic ground state extrapolates to $(0.28 \\pm 0.08)$~erg~s$^{-1}$cm$^{-2}$sr$^{-1}$.  Over the ISO-SWS aperture this amounts to $(17 \\pm 5)~\\rm L_{\\sun}$.  Compared with the total observed $\\rm H_2$ line emission (after extinction correction) of $(0.16 \\pm 0.05)$~erg~s$^{-1}$cm$^{-2}$sr$^{-1}$, we find that our line spectra account for more than half of the total $\\rm H_2$ emission.  Our observations target the brightest field in the Orion outflow. The outflow covers an area of about $2\\arcmin \\times 2\\arcmin$.  The average H$_2$ brightness over this area we estimate from the 1-0~S(1) map of Garden (\\cite{gar86}) to approximately 20\\% of that in our observed field, so that the total $\\rm H_2$ luminosity of the OMC-1 outflow is estimated to be $(120 \\pm 60)~ \\rm L_{\\sun}$.  This is consistent with the 94~$\\rm L_{\\sun}$ estimated by Burton \\& Puxley (\\cite{bur90b}).   \\subsubsection{What excites the highest-energy levels?} \\label{extraexcitation}  Table~\\ref{column} and Fig.~\\ref{temp} illustrate that only a small fraction of the warm molecular gas is at the high excitation temperatures, which reach 3000~K. This is difficult to reconcile with the expected smooth temperature profile of a single planar C-type shock, in which the gas temperature changes smoothly, and where a large fraction of the warm gas is near the maximum temperature (Timmermann \\cite{tim96b}). Even with a distribution of shock speeds, and a correspondingly wide range in peak temperatures, an excitation temperature distribution similar to that shown in Fig.~\\ref{temp} is difficult to understand. It would require a velocity distribution where only a small fraction, about 1\\%, of the gas is shocked at the high speed necessary to produce a 3000 K excitation. In bow shocks, e.g., the velocity changes slowly with distance from the apex, and such a distribution of velocities would not be expected.  In dissociative J-type shocks, the molecules are destroyed in the shock, and they reform in a postshock layer where the temperature has dropped much below 3000 K, somewhat dependent on the H$_2$ formation rate efficiency at higher temperatures, which is essentially unknown (e.g. Bertoldi \\cite{ber97}).  Dissociative J-shocks can therefore not account for the high excitation H$_2$ we observe.  Even if temperatures of 3000 K or more can be reached in non-dissociative shocks, the higher H$_2$ levels would remain subthermally excited unless the gas density is high enough that the collisional excitation and deexcitation rates are comparable to those for radiative decay. A ``critical'' gas density can be defined for a given level as that for which the total collisional deexcitation rate of this level equals its total radiative decay rate.  In Fig.~\\ref{ncrit} we plot the critical density computed this way for states in the vibrational ground state, $v=0$, up to $J=16$. We see that even at kinetic temperatures of 3000 K, gas densities above $10^6\\rm cm^{-3}$ would be necessary to maintain the high $v=0$ levels at populations resembling LTE.  Since such high densities may not prevail in the shocked gas of the Orion outflow, we may explore mechanisms other than thermal excitation that could account for the population of the higher energy states (see also Bertoldi et al.~\\cite{ber00}).  \\begin{figure}[htb] \\begin{center} \\includegraphics[width=1.0\\columnwidth]{figure10.ps}  % \\caption{ Critical densities, $n_{\\rm crit}$, of $v=0$ states of H$_2$ as a function of level energy, for H$_2$--H$_2$ and H$_2$--H collisions. The critical density is the ratio between the sum of all radiative rate (Einstein-$A$) coefficients, and the sum of the collisional deexcitation rate coefficients from a state to all states with lower energy. Einstein-$A$ coefficients were adopted from Turner et al.~(1977), and collisional rate coefficients from Bourlot et al.~(1999).} \\label{ncrit} \\end{center} \\end{figure}  \\paragraph{Time-dependent C-shocks:} When a high velocity outflow strikes dense molecular gas and thereby a C-type shock is first established, J-type shocks can temporarily form within the C-shock. In such an embedded J-shock, a small column of gas is heated to high temperatures (Chi\\`{e}ze et al. \\cite{chi98}), and if the density is sufficiently high, this could account for the high-excitation tail of the column density distribution (Flower \\& Pineau des For\\^ets \\cite{flo99}).  The lifetime of the embedded J-shock is small, so that the high-temperature excitation tail would be a transient phenomenon, unless shocks are constantly reforming. Embedded J-shocks may also form when a C-shock encounters dense clumps.  \\paragraph{Formation pumping:} Another possible pumping mechanism of the high-energy states is the formation of H$_2$. Molecular hydrogen is believed to form on the surfaces of dust grains.  Some of the 4.5~eV released during the formation of an H$_2$ molecule is used up to leave the grain, and the remainder is split between translation, rotation, and vibration of the new molecule. The exact level distribution of newly formed H$_2$ is yet unknown, but it could very well contribute to the observed excitation at intermediate energies, $E\\approx 1-3$~eV (Black \\& van Dishoeck \\cite{bla87}; Le Bourlot et al. \\cite{leb95}).  Using Eq.~\\ref{n_tot} we can sum up the column densities of all levels with energy $E/k \\geq 10 \\,000$~K, to find a column density $1.30 \\times 10^{18}$~cm$^{-2}$, a fraction $6.8 \\times 10^{-4}$ of the total warm H$_2$ column.  Could H$_2$ formation in a steady state produce such a fraction of molecules in highly excited states?  The pumping rate due to formation pumping is equal to the H$_2$ formation rate, $n({\\rm H}) n_{\\rm H} R_{\\rm gr}$, where $R_{\\rm gr} \\approx 5\\times 10^{-17}$~cm$^3$~s$^{-1}$ is the H$_2$ formation rate coefficient per hydrogen nucleus. We estimate the radiative decay rate by starting with the characteristic radiative lifetime of $\\sim 10^6$~sec for a molecule in a vibrational level $v \\approx 5$, and note that $\\sim 5$ jumps may be required to reach the ground state, so that the effective $A$-coefficient $A_{\\rm x} \\approx 2 \\times 10^{-7}~s^{-1}$. The population balance for the excited states then writes \\begin{equation} R_{\\rm gr} ~n_{\\rm H}~ n({\\rm H}) ~=~ n_{\\rm x}({\\rm H_2}) ~A_{\\rm x}, \\end{equation} which yields an excited H$_2$ fraction \\begin{eqnarray} \\label{eq:excited} {n_{\\rm x}({\\rm H_2})\\over n({\\rm H_2})} & = & {n{(\\rm H)}\\over n({\\rm H_2})} {n_{\\rm H} R_{\\rm gr} \\over A_{\\rm x} } \\nonumber \\\\ & = & 5\\times 10^{-4} \\left( {n_{\\rm H}\\over 10^6{\\rm cm^{-3}}}~~ {n{(\\rm H)}\\over 2n({\\rm H_2})} \\right)~, \\end{eqnarray} which would be consistent with the observed value, if the term in brackets assumes a value of order unity.  This simple estimate thus shows that H$_2$ formation could account for some of the high excitation level populations if the density is high, the atomic fraction not small, and the formation rate coefficient in the warm shocked gas is somewhat higher than the value implied at $\\approx 100$~K from Copernicus observations, which is $R_{\\rm gr} \\approx 3\\times 10^{-17}$~cm$^3$~s$^{-1}$ (Jura 1975).  To illustrate the possible importance of H$_2$ formation for the high-excitation level pumping we show that a simple superposition of two gas layers with hydrogen nuclei density $n_{\\rm H}=10^6\\rm cm^{-3}$, atomic fraction $n({\\rm H})/n_{\\rm H}=0.5$, and respective temperatures of 200 K and 800 K, with column densities $N_{\\rm H_2}=1.2\\times 10^{22}\\rm cm^{-2}$ and $1.2\\times 10^{21}\\rm cm^{-2}$, can in fact reproduce the observed level column distribution better than any shock model currently available.  We used the photodissociation front code of Draine \\& Bertoldi (1996), but without UV illumination, to compute the non-LTE level distributions for gas at a fixed temperature, density, and molecular fraction. We include H$_2$ formation with a rate coefficient $R_{\\rm gr}~=~5\\times 10^{-17} \\rm cm^{3} s^{-1}$ and assume a level distribution for the newly formed H$_2$ following $N(v,J) \\propto (2J+1) e^{-E(v,J)/kT_{form}}$, with a ``formation temperature'' $T_{form}=5000$~K chosen to match the slope of the observed high-excitation level distribution.  Fig.~\\ref{fig:twoT} illustrates how the newly-formed H$_2$ molecules give rise to a high-excitation tail in the levels' column density distribution.  The remaining gas displays a thermal distribution at least up to the levels which are mainly populated by H$_2$ formation.  \\begin{figure}[htb] \\begin{center} \\includegraphics[width=1.05\\columnwidth]{figure11.ps} \\caption{ Level column density distributions for two gas layers at temperatures 200 K and 800 K with H$_2$ column densities of $1.2\\times 10^{22}\\rm cm^{-3}$ and $1.2\\times 10^{21}\\rm cm^{-3}$, respectively.  Individual vibrational level distributions are shown as separate lines.  The sum of both distributions well matches the observed level distribution toward Peak 1 which is shown in Fig. 7.} \\label{fig:twoT} \\end{center} \\end{figure}  \\paragraph{Non-thermal collisions:} An even more important pumping mechanism for the high-excitation levels may be non-thermal collisions between molecules and ions in a magnetic shock.  In magnetic C-type shocks, which are believed to be responsible for most of the emission in Peak~1, the gas is accelerated through fast inelastic collisions. In a magnetic precursor the ions, which are tied to the magnetic field, collide with the undisturbed pre-shock gas at relative velocities comparable to the shock speed. Such non-thermal ion--molecule collisions lead to the acceleration of the molecules and to their internal excitation. High-velocity molecules subsequently collide with other molecules, leading to a cascade of collisions during which the relative kinetic energy is in part converted to internal excitation of the molecules (O'Brien \\& Drury \\cite{bri96}).  In sufficiently fast C-shocks, the ion--H$_2$ and H$_2$--H$_2$ collisions can even lead to a significant collisional dissociation rate.  The molecules dissociated in a steady-state, partially dissociative shock reform further downstream, so that across such a shock the H$_2$ dissociation rate equals the H$_2$ reformation rate. For every collisionally dissociated molecule there will be a larger number of inelastic collisions which did not lead to dissociation, but to the excitation of the molecule into high ro-vibrational states, up to the dissociation limit.  The high-excitation H$_2$ level column densities thereby created should therefore be larger than those caused by H$_2$ formation alone.  Note that such energetic collision between ions and H$_2$ in C-shocks are relatively infrequent because ions are rare -- thus the excited H$_2$ has time enough to cascade to lower levels between collisions, giving rise to line emission from the highly excited levels. In C-type shocks, however, dissociations take place too quickly for highly excited H$_2$ to radiatively decay.  We conclude that non-thermal collisions in partially dissociative C-shocks could pump the high-excitation states in the H$_2$ electronic ground state to the levels observed. However, no detailed shock models are available yet which account for this process.  \\paragraph{0 -- 0 S(25):} The $J=27$ level observed through the $0-0$~S(25) line appears overpopulated by a factor of seven over what would be expected from the least-squares fit of the data shown in Fig.~\\ref{excit}.  The $J=27$ level is 3.6~eV above ground and only 0.9~eV from the dissociation limit. H$_2$ molecules which are newly formed on grains are unlikely to be able to populate states so high, because some fraction of the formation energy is lost to overcome the grain surface potential, and some goes to kinetic and vibrational excitation.  Unless we misidentified the $0-0$~S(25) line, it appears that a different mechanism may be populating this level and possibly other high levels. The gas-phase formation of H$_2$ via H$^-$, e.g., might be able to leave the new molecule in such a high rotational state (Bieniek \\& Dalgarno \\cite{bie79}; Black et al. \\cite{bla81}; Launay et al. \\cite{lau91}).  \\subsubsection{Comparison with shock models} \\label{model}  For over 20 years, evidence accumulated that the $\\rm H_2$ emission from OMC-1 may arise from shocks (Gautier et al.  \\cite{gau76}; Kwan \\& Scoville \\cite{kwa76}).  However, the physical nature of these shocks remains unclear. Models for planar J-type (Hollenbach \\& Shull \\cite{hol77}; Kwan \\cite{kwa77}; London et al.  \\cite{lon77}) or C-type (Draine \\cite{dra80}; Draine \\& Roberge \\cite{dra82}; Chernoff et al. \\cite{che82}; Draine et al. \\cite{dra83}) shocks were unable to reproduce the observed wide velocity profiles (Nadeau \\& Geballe \\cite{nad79}; Brand et al.  \\cite{bra89b}; Moorhouse et al. \\cite{moo90}; Chrysostomou et al.  \\cite{chr97}), or the wide range of excitation conditions observed.  Bow shocks were suggested to account for the observed range of excitation conditions and the wide velocity profiles (Hartigan et al.  \\cite{har87}; e.~g. Smith et al. \\cite{smi91a,smi91b}), but it remains unclear whether these are predominantly C-type, J-type, or a combination of those.  \\paragraph{Planar shock models:} In Sect.~\\ref{finestructure} we suggested that some fraction of the fine structure line emission may arise from dissociative J-shocks with velocities of about 85~km~s$^{-1}$, pre-shock densities $n_{\\rm H}\\approx \\rm 10^5 - 10^6\\,cm^{-3}$, and a beam filling factor of 3--4 (Hollenbach \\& McKee \\cite{hol89}).  In such dissociative shocks most of the excitation of the $v=0$, $ J \\leq 5$ levels is collisional, and the emission arises in the H$_2$ reformation region where the temperature levels at 400 to 500~K, which is somewhat below the observed excitation temperature of the lowest levels.  The higher levels would then be predominantly pumped by newly formed molecules.  Such a model however neither fits the low excitation nor the higher excitation level populations very well.  The deficits of the J-shock model could be compensated if we combined it with a C-shock model, e.g., one of Kaufman \\& Neufeld (\\cite{kau96}) with $v_{\\rm s}=25$~$\\rm km~s^{-1}$, $n_{\\rm H_2}=10^5-10^6~{\\rm cm^{-3}}$, and a beam filling factor 0.3. Such a combined model provides a good fit to the $v=0$, $J=3$ to 9 level populations, although higher rotational level populations are predicted too large  (see Fig.~\\ref{ex_model}: HK).  A combination of J-type and C-type shocks would be consistent with the picture proposed by Chernoff et al.~(\\cite{che82}), who suggested that a high velocity, $\\sim {\\rm 110~km~s^{-1}}$, wind emanating from an object near IRc2 drives a $\\sim {\\rm 30~km~s^{-1}}$ expanding shell of swept-up material.  The low beam filling factor of the C-shock emission could be due to the clumping of the ambient medium.  \\begin{figure}[htb] \\begin{center} \\includegraphics[width=1.\\columnwidth]{figure12.ps} \\caption{ Comparison of the observed H$_2$ level column distribution (solid line) with models: dissociative J-shock plus C-shock model (HK: long-dashed), J-type cooling flow model (B88: dash-dotted), bow shock model (S91: short-dashed), combination of two planar C-shock models (KN96: dotted).} \\label{ex_model} \\end{center} \\end{figure}  The $\\rm H_2$ level populations implied by previous ground-based observations (e.g. Brand et al. \\cite{bra88}; Parmar et al. \\cite{par94}; Burton \\& Haas \\cite{bur97}) were attempted to match with an empirical planar J-shock ``cooling flow\" model (Brand et al. \\cite{bra88}; Chang \\& Martin \\cite{cha91}; Burton \\& Haas \\cite{bur97}), which assumes that the cooling is dominated by $\\rm H_2$, and cooling by other molecules such as H$_2$O and CO may be neglected. Such a model can match the medium and high-excitation level populations, although it somewhat overestimates the population of lower rotational levels. These models assume LTE level distributions, which as we argued above, may not be a valid assumption for the high-excitation levels if the gas density is below $10^6\\rm cm^{-3}$.  Furthermore, theoretical chemical studies show that most oxygen not locked in CO is converted to $\\rm H_2O$ (Draine et al. \\cite{dra83}; Kaufman \\& Neufeld \\cite{kau96}), which is detected by ISO-LWS observations toward OMC-1 (Harwit et al. \\cite{har98}; Cernicharo et al. \\cite{cer99}). Both H$_2$O and CO should therefore be significant coolants, and the neglect of this in these models is worrisome.  Currently available more realistic single shock models do not seem to fit the observed H$_2$ level distribution.  It appears necessary to combine at least two shock models, one to account for the high-excitation level populations, one for the low-excitation levels. For example, combining two models from Kaufman \\& Neufeld (\\cite{kau96}), with shock velocities of 20 and 40 km~s$^{-1}$, and beam filling factors of 1 and 0.026, respectively, can well match the level population up to $E/k\\approx 20\\,000$ K (Fig.~\\ref{excit}, KN96); a pre-shock H$_2$ number density of $\\rm 3~\\times~10^{5}~cm^{-3}$ was adopted. The Kaufman \\& Neufeld models however do not account for time-dependency, formation pumping, or non-thermal excitation.  \\paragraph{H$_2$ velocity dispersion:} Optical and near-IR observations with high spectral resolution toward the OMC-1 outflow show that typical FWHM widths of the H$_2$ lines are 50--60~km s$^{-1}$ (Nadeau \\& Geballe \\cite{nad79}; Moorhouse et al. \\cite{moo90}; Geballe \\& Garden \\cite{geb87}; Chrysostomou et al. \\cite{chr97}), and that the line wings can extend to several hundred km s$^{-1}$ (Ramsey-Howat et al., in prep.). Molecular hydrogen is expected to be destroyed in shocks with velocities larger than 30 to 50 km s$^{-1}$, depending on the magnetic field strength.  It is therefore puzzeling how the H$_2$ emission can show such large velocity dispersions, even in filaments which are only several arcseconds in size.  \\paragraph{Bow-shocks:} It has been suggested that the H$_2$ emission arises in bow-shocks, in which the effective shock velocity decreases from the tip to the wake. The shock speed at the apex may be high enough to produce a dissociative J-shock here. But further down the wake, non-dissociative C-type shocks can prevail, with peak temperatures in the shocked molecular layers that decrease steadily down the wake.  Thereby a large range of temperatures for the molecular gas exists in a single bow shock. This could account for the observed level excitation, and may also explain the observed constancy of $\\rm H_2$ excitation over the entire OMC-1 outflow (Brand et al.  \\cite{bra89a}).  The existence of bow-shocks is also supported by the observation of double-peaked velocity profiles for isolated regions in the outflow (Chrysostomou et al. \\cite{chr97}), and of knots of [Fe~{\\sc ii}] emission which coincide with ``fingers'' of $\\rm H_2$ emission.  Allen \\& Burton (\\cite{all93}) suggest that the [Fe~{\\sc ii}] and $\\rm H_2$ emission trace tips and wakes of bow shocks formed in a stellar outflow.  Recent observations of [Fe~{\\sc ii}] and $\\rm H_2$~1-0~S(1) velocity profiles (Tedds et al. \\cite{ted99}) however question the bow shock picture.  Alternatively, Stone et al. (\\cite{sto95b}) proposed that the bows result from Rayleigh-Taylor instabilities when a poorly collimated outflow accelerates in an ambient medium of decreasing density, or when catching up with a slower shock.  It is difficult to understand how the high velocity excited H$_2$ can be produced in a bow shock which is produced, e.g., by a dense bullet which is moving through a medium initially at rest. Ambient gas which has passed through parts of the bow shock which are not strong enough to dissociate the H$_2$ will not be accelerated to velocities much larger than 30 km s$^{-1}$, unless the magnetic field is very strong. If alternatively the bow shock arises from a molecular wind impinging on a dense obstacle, then the problem arises how the molecular wind was accelerated to over 100 km s$^{-1}$ without destroying the molecules, and why we do not see a lot more mass, traced by CO, e.g., at such high velocities.  Smith et al.  (\\cite{smi91a}) are able to reproduce the shape and width of the observed H$_2$ lines in the Orion outflow with bow shock models, but only by assuming a magnetic field strength of 50 mG, significantly higher than the 10 mG implied by polarization studies (Chrysostomou et al. \\cite{chr94}).  In Fig.~\\ref{ex_model}, we compare our data with a bow shock model by Smith (\\cite{smi91b}), adopting a peak shock velocity of ${\\rm 100~km~s^{-1}}$ and an Alfv\\'en speed of ${\\rm 2~km~s^{-1}}$.  This model is able to match the medium-excitation level populations well, but underestimates the low-excitation, and overestimates the high-excitation levels.  Note that the H$_2$ excitation in the models of Smith, like for the planar J-shock model of Brand et al. (\\cite{bra88}), was calculated under the assumption of LTE, and also it ignores non-thermal excitation mechanisms. These models therefore overestimate the population of the high energy levels by thermal collision, and at the same time they underestimate the level population because they neglect non-thermal excitation mechanisms.  We conclude that current shock models are able to reproduce the overall H$_2$ level distribution only when combining shocks with a range of velocities. However, most models do not include the physics most likely to account for the highest excitation level populations."
    },
    "0002/astro-ph0002330_arXiv.txt": {
        "abstract": "We examine the effect of  primordial dark matter velocity dispersion and/or particle self-interactions on the structure and stability of galaxy halos, especially with respect to the formation of  substructure and central density cusps. Primordial velocity dispersion  is characterised by a ``phase density'' $Q\\equiv \\rho/\\langle v^2\\rangle^{3/2}$,  which for relativistically-decoupled relics is determined by   particle mass and spin and is  insensitive to cosmological parameters. Finite $Q$ leads  to small-scale  filtering    of  the primordial  power spectrum, which reduces substructure, and limits the maximum central density of halos, which eliminates central cusps. The relationship between $Q$ and halo observables is estimated.  The primordial $Q$ may be preserved in the cores of halos and if so leads to a predicted relation, closely analogous to that in degenerate dwarf stars, between the central density  and  velocity dispersion.  Classical polytrope solutions are used to model  the structure of  halos of collisional dark matter, and to show that self-interactions in halos today are probably not significant because they   destabilize halo cores via heat conduction.  Constraints on masses and self-interactions of dark matter particles are estimated from halo stability and other considerations. ",
        "introduction": "The successful concordance  of predictions and observations of large scale structure and microwave anisotropy vindicates many assumptions of standard cosmology, in particular the hypothesis that the dark matter is composed of primordial particles which are cold and collisionless\\cite{cdm}. At the same time, there are hints of discrepancies observed in    the small-scale structure within galaxy haloes, which we explore as two related but separate issues, namely the predictions of excessive substructure and sharp central cusps in dark matter halos.  The first ``substructure problem'' is that    CDM  predicts excessive relic substructure\\cite{ghigna,klypin}: much of the mass of a CDM halo is not smoothly distributed but is concentrated in many massive sublumps,   like galaxies in galaxy clusters. The model predicts that galaxy halos should contain many dwarf galaxies which  are not seen, and which would disrupt disks even if they are invisible. The substructure problem  appears to be  caused by the ``bottom-up'' hierarchical clustering predicted by CDM power spectra; fluctuations on small scales collapse early and survive as dense condensations. Its absence  hints that the small scale power spectrum is filtered  to suppress early collapse on subgalactic scales.  The second ``cusp problem'' is that CDM  also predicts\\cite{dubinski,nfw,moore98,moore99a,moore99b} a universal, monotonic increase of  density towards the center of  halos  which is not seen  in close studies of dark-matter-dominated galaxies\\cite{swaters,swaters2,carignan} (although the observational issue is far from settled\\cite{vandenbosch,dalcanton}). The formation of central cusps has been observed for many years in simulations of collapse of cold matter in a wide variety of circumstances; it may be thought of as   low-entropy material sinking to the center during halo formation. Simulations suggest that dynamical ``pre-heating'' of CDM by hierarchical clustering is not enough to prevent a cusp from forming--- that some material is always left with a low entropy and sinks to the center. If this is right, the central structure of halos might provide clues to the primordial entropy which are insensitive to complicated details of nonlinear collapse.    It may be  possible to explain these discrepancies in a CDM framework\\cite{frank}, for example by using various baryonic contrivances. It is also possible that the observations can be interpreted more sympathetically for CDM; we explore this possibility in more detail in a separate paper\\cite{dalcanton}. However it is also possible that the problems with halo structure are giving specific quantitative clues about new properties of the dark matter particles. By examining halo structure and stability, in this paper we    make a quantitative assessment of the effects   of   modifications of the two main properties of CDM--- the addition of primordial velocity dispersion, and/or the addition of particle self-interactions. In particular we focus on aspects of halo structure which provide the cleanest ``laboratories'' for studying dark matter properties.  The ultimate goal of this exercise is  to measure and constrain particle properties from halo structure.   Endowing the particles with   non-zero primordial velocity dispersion produces two separate effects: a filter in the primordial power spectrum which limits small-scale substructure, and a  phase packing  or Liouville limit  which produces halo cores. Both effects depend on the same   quantity, the ``phase density'' which we choose to define using the most observationally accessible units, $Q\\equiv \\rho/\\langle v^2\\rangle^{3/2}$, where $\\rho$ is the density and $\\langle v^2\\rangle$ is the velocity dispersion. The definitions of these quantities depend on whether we are discussing fine-grained or coarse-grained $Q$.\\footnote{For a uniform monatomic ideal thermal gas, $Q$ is related in a straightforward way to the  usual  thermodynamic entropy; for  $N$ particles, $S=- kN[\\ln (Q)+{\\rm constant}].$} For collisionless particles, the fine-grained $Q$ does not change but the coarse-grained $Q$ can decrease as the   sheet occupied by particles folds up in phase space. The coarse-grained $Q$ can be estimated directly from astronomical observations, while the fine-grained $Q$ relates directly to microphysics of dark matter particles. For particles which decouple when still relativistic, the initial microscopic phase density  $Q_0$, which for nondissipationless collisionless particles is the maximum value for all time,  can be related to the particle mass and type, with little reference to the cosmology. The most familiar examples are the standard neutrinos, but  we include in our discussion the more general case which yields different numerical factors for bosons and for particles with a significant chemical potential.  The physics of both  the filtering and the phase packing in the collisionless case  closely parallels that of massive neutrinos\\cite{gerstein,cowsik},  the standard form of ``hot'' dark matter. Dominant hot dark matter  overdoes both of these effects--- the filtering scale  is too large to agree with observations of galaxy formation (both in   emission and quasar absorption) and the phase density is too low to agree with observations of giant-galaxy halos\\cite{tremaine}.  However one can introduce new particles with a lower velocity dispersion (``warm'' dark matter, \\cite{bond,bardeen,blumenthal,melott,primack,dodelson}), which is the option we consider here.  Although warm dark matter has most often  been invoked as a solution to fixing apparent (and no longer problematic\\cite{peacock,cdm}) difficulties with predictions of the CDM power spectrum for matching galaxy clustering data, a spectrum filtered on smaller scales may also solve  several other classic problems of CDM on galactic and subgalactic scales\\cite{whiterees,navarrosteinmetz} which are sometimes attributed to baryonic effects. The  main effect in warm models is that the first nonlinear objects are   larger and form   later, suppressing substructure and increasing the angular momentum of galaxies\\cite{sommerlarsen}. This improves the predictions for dwarf galaxy populations\\cite{colin}, baryon-to-dark-matter ratio, disk  size and angular momentum, and quiet flows on the scale of galaxy groups. If the filtering is confined to small scales the predictions are likely to remain acceptable for Lyman-$\\alpha$ absorption during the epoch of galaxy formation at $z\\approx 3$\\cite{croft,dave,white}.  Liouville's theorem tells us that dissipationless, collisionless particles can only decrease their coarse-grained phase density, and we conjecture that  halo   cores on small scales approximately preserve the primordial phase density. The universal character of the phase density allows us to make definite predictions for the scaling of core density and  core radius with halo velocity dispersion. These relations are analogous to those governing nonrelativistic degenerate-dwarf stars: more tightly bound (i.e. massive) halos should have smaller, denser cores. A  survey\\cite{dalcanton} of available evidence on the phase density of dark matter cores on scales from dwarf spheroidal galaxies to galaxy clusters shows that the   phase density needed to create the cores of rotating dwarf  galaxies is much lower than that apparently present in dwarf spheroidal galaxies\\cite{aaronson,olszewski,faber,lake,gerhard,ralston,mateo}--- so at  least one of these populations is not probing primordial phase density. Translating into masses of neutrino-like relics, the spheroidals prefer masses of about 1 keV (unless the observed stars occupy only a small central portion of an implausibly large, massive and high-dispersion halo), and the disks prefer about 200 eV.  The larger phase density is also preferred from the point of view of filtering.  If we take $\\Omega\\approx 0.3$ (instead of 1  as in most of the original warm scenarios--- which reduces the scale for a given mass, because it lowers the temperature and therefore the number of the particles), the filtering scale for 1 keV particles is at about $k=3{\\rm Mpc}^{-1}$--- small enough to preserve the successful large-scale predictions of CDM but also large enough to impact the substructure problem. Galaxy halo substructure therefore favors a primordial phase density corresponding to collisionless thermal relics with a mass of around 1 keV. In this scenario the densest dwarf spheroidals might well preserve the primordial phase density and in principle could allow a measurement of the particle mass.\\footnote{This raises another unresolved issue: whether the filtering actually prevents systems as small as dwarf spheroidals from forming at all. The predictions of warm dark models are not yet worked out enough to answer this question.} (Conversely, a  mass as large as 1 keV can only solve the core problem in disks with additional  nonlinear dynamical heating, so that the central matter  no longer remains on the lowest adiabat, or with the aid of baryonic effects\\cite{frank}.)  To have the right mean density and phase density today,  relativistically-decoupling particles of this phase density  must have separated out at least as early as the QCD era, when the number of degrees of freedom was much larger than at classical weak decoupling. Their interactions with normal Standard Model particles must therefore be ``weaker than weak,'' ruling out not only standard neutrinos but many other particle candidates. The leading CDM particle candidates, such as   WIMPs and axions,   form in standard scenarios with much higher  phase densities, although more elaborate mechanisms are possible to endow these particles with the  velocities to dilute $Q$. We   review briefly some of the available options for making low-$Q$ candidates, such as particles decaying out of equilibrium.  A new wrinkle on this story comes if we  endow the particles with self-interactions\\cite{carlson,delaix,atrio,spergel,mohapatra}. We  consider a simple parametrized model of particle self-interactions based on massive intermediate particles of adjustable mass and coupling, and explore the constraints on these parameters from halo structure. Self-interactions change the filtering of the power spectrum early on, and if they are strong enough they qualititatively change the global structure and stability of halos.  In the interacting case, linear perturbations below the Jeans scale   oscillate as sound waves instead of damping by free streaming--- analogous to a baryon plasma rather than a neutrino gas. This effect  introduces a filter which is sharper in $k$ than that from streaming, and also on a scale about ten times smaller than the streaming for the same rms particle velocity--- about right to reconcile the appropriate filtering scale with the $Q$ needed for phase-density-limited disk cores. These self-interactions could be so weak that the particles are   effectively collisionless today as in standard CDM.  On the other hand stronger self-interactions have major  effects during the nonlinear stages of  structure formation and on the structure of galaxy halos\\cite{spergel}. We consider this possibility in some detail, using Lane-Emden polytropes as   fiducial models for collisional halos. Their structures are close analogs of degenerate dwarf stars and we call them ``giant dwarfs''.  We find that these structures are subject to an instability caused by heat conduction by particle diffusion.\\footnote{Degenerate dwarf stars  are not subject to this instability because they are supported without a temperature gradient; the same stabilization could occur in halo cores only if the dark matter is fermionic and degenerate (e.g., \\cite{fuller,shi}). The instability we discuss here is  essentially what   happens in a thermally-supported star with no nuclear reactions, except that the conduction is by particle diffusion rather than by radiation. This effect   may have already been observed numerically.\\cite{hannestad}} Although a little of this might be interesting (e.g. leading to the formation of central black holes\\cite{ostriker} or to high-density, dwarf spheroidal galaxies),   typical halos can only be significantly collisional if they last for a Hubble time;  for this to be the case,  the particle interactions must be so strong that diffusion is suppressed, which in turn requires a fluid behavior for all bound dark matter structures. This option is not very attractive from a phenomenological point of view\\cite{delaix,spergel}; for example, dwarf galaxies or galaxies in clusters tend to sink like rocks instead of orbiting like satellites, and the collapse of cores occurs most easily in those low-dispersion halos where we seek to stabilize them. ",
        "conclusions": "We have found that some halos might preserve in their inner structure observable  clues to new dark matter physics, and   that  indeed some current observations already hint that the dark matter might be warm rather than cold. We conclude with a summary:  1. Halo cores can be created by a ``phase-packing limit'' depending on finite  initial phase density. They may provide a direct probe of primordial velocity dispersion in dissipationless dark matter.  2. For relativistically-decoupled thermal relics, the phase density depends on the particle mass and spin but not on cosmological parameters.   3. Rotation curves in a few dwarf disk  galaxies indicate cores with a phase density  corresponding to that  of a 200 eV thermal relic or  an rms velocity of about 0.4 km/sec at the current cosmic mean density. Velocity dispersions in dwarf spheroidal galaxies indicate a  higher phase density, corresponding to a thermal relic mass of about 1 keV. At most one of these populations can be tracing the primordial phase density.  4. Thermal relics in this mass range can match the mean cosmic density with a plausible superweak decoupling from Standard Model particles before the QCD epoch.  5. Other very different particles are consistent with the halo data, provided they have the about the same mean density and phase density. Examples include    WIMPs from particle decay and axions from defect decay.  6. Cores due to phase packing limited by primoridial $Q_0$ predict a universal relation between core radius and halo velocity dispersion. The relation is not found in a straightforward interpretation of the data.  7. Primordial velocity dispersion  also suppresses halo substructure (and solves some other difficulties with CDM)  by filtering primordial adiabatic perturbations. Estimates based on luminosity functions prefer filtering  on a scale of about $k\\approx 3 {\\rm Mpc}^{-1}$; for collisionless particles, this scale corresponds to a filter caused by streaming of  about a 1keV thermal relic.  8. Weak self-interactions change  from streaming to acoustic behavior, reducing  the damping  scale and sharpening the filter.  9. Stronger self-interactions destabilize halos by thermal conduction, making the cusp problem worse (unless they are very strong--- too strong for satellite-galaxy  kinematics--- or particles are degenerate, eliminating the central temperature gradient).   10. A simulation which samples a warm distribution function reasonably well is strongly motivated, to determine whether primordial $Q$ is preserved in the centers of halos, or whether nonlinear effects can amplify dynamical heating in such models to explain cores on all scales."
    },
    "0002/astro-ph0002106_arXiv.txt": {
        "abstract": "The radiative cooling of optically thin gaseous regions and the formation of a two-phase medium and of cold gas clouds with a clumpy substructure is investigated.  We demonstrate how clumpiness can emerge as a result of thermal instability. In optically thin clouds, the growth rate of small density perturbations is independent of their length scale as long as the perturbations can adjust to an isobaric state. However, the growth of a perturbation is limited by its transition from isobaric to isochoric cooling when the cooling time scale is reduced below the sound crossing time scale across its length scale. The temperature at which this transition occurs decreases with the length scale of the perturbation. Consequently small scale perturbations have the potential to reach higher amplitudes than large scale perturbations.  When the amplitude becomes nonlinear, advection overtakes the pressure gradient in promoting the compression  resulting in an accelerated growth of the disturbance.  The critical temperature for transition depends on the initial amplitude.  The fluctuations which can first reach nonlinearity before their isobaric to isochoric transition will determine the characteristic size and mass of the cold dense clumps which would emerge from the cooling of an initially nearly homogeneous region of gas. Thermal conduction is in general very efficient in erasing isobaric, small-scale fluctuations, suppressing a cooling instability. A weak, tangled magnetic field can however reduce the conductive heat flux enough for low-amplitude fluctuations to grow isobarically and become non-linear if their length scales are of order $10^{-2}$ pc. If the amplitude of the initial perturbations is a decreasing function of the wavelength, the size of the emerging clumps will decrease with increasing magnetic field strength. Finally, we demonstrate how a 2-phase medium, with cold clumps being pressure confined in a diffuse hot residual background component, would be sustained if there is adequate heating to compensate the energy loss. ",
        "introduction": "The interstellar medium and cold gas clouds are characterized by a clumpy substructure and a turbulent velocity field (Larson 1981, Blitz 1993).  As molecular clouds are the sites of star formation, their formation, internal structure and dynamics determines the rate of star formation and the properties of young stars, such as their mass function or binarity.  The understanding of the origin of cold clouds and their internal substructure is therefore of fundamental importance for a consistent theory of star formation and galactic evolution.  In the nearby clouds, the dispersion velocity inferred from molecular line width is often larger than the gas sound speed inferred from the line transition temperatures (Solomon et al 1987). MHD turbulence may be responsible for the stirring of these clouds (Arons \\& Max 1975).  This conjecture is supported by the polarization maps and direct measurements of field strength in some star forming regions (Myers \\& Goodman 1988, Crutcher et al. 1993). Recent simulations of MHD turbulence, however, suggest that it dissipates rapidly (Gammie \\& Ostriker 1996, MacLow et al 1998, MacLow 1999, Ostriker et al. 1999).  One possible source of energy supply is winds and outflows from young stellar objects (Franco \\& Cox 1983, McKee 1989).  But in regions where star formation is inactive, clumpy structure with velocity dispersion is also observed.  Thus, the origin and energy supply of clumpy cloud structure remains an outstanding issue.  On small scales, magnetic field pressure is important in regulating infall and collapse of protostellar clouds and the formation of low-mass stars (Mouschovias \\& Spitzer 1976, Nakano 1979, Shu 1993). For clouds with sub-critical masses, gravitational contraction is proceeded by ambipolar diffusion which for typical cloud densities operates on a timescale $\\tau_B \\sim 10^{7-8}$ yr (Lizano \\& Shu 1989, Mouschovias 1991). In regions with intense star formation activities such as the central region of Orion, $\\tau_B$ for individual dense clumps is comparable to the typical age of the young stellar objects.  But, the spread in stellar ages ($\\Delta \\tau_\\ast \\sim 10^6$ yrs) appears to be considerably shorter than $\\tau_B$ (Carpenter et al. 1997, Hillenbrand, 1997). This coeval star formation history requires either a coordinated trigger mechanism for star formation within initially magnetically supported clumps or subcritical collapse, fragmentation and star formation of a larger molecular cloud region in which the magnetic field plays a weak role.  A rapid and coordinated episode of star formation can also be inferred in globular clusters (Brown et al. 1991, 1995, Murray \\& Lin 1992, Lin \\& Murray 1992).  In some metal deficient clusters such as M92, the total amount of heavy elements corresponds to the yield of a few supernovae.  If star formation has proceeded over a duration $\\Delta \\tau_\\ast$ comparable to the expected life span ($\\sim$ a few $10^6$ yrs) of massive stars, a significant metallicity spread would be expected, in contrast to the observations (e.g. Kraft 1979). At least in these systems, $\\Delta \\tau_\\ast <\\tau_B$ and star formation may have proceeded through supercritical collapse.  The dynamical timescale of most clusters at their half mass radius is $\\tau_d \\approx 10^6$ yr.  Any energy dissipation associated with the episode of star formation would imply an even longer dynamical timescale in the proto cluster cloud prior to that event.  We infer that $\\Delta \\tau_{\\ast}$ was comparable to or shorter than the dynamical timescale of the proto cluster clouds indicating a rapid fragmentation and star formation episode.  In this paper, we focus on the rapid emergence of clumpy structure during the formation and collapse of a thermally unstable supercritical cloud. This process is relevant to the formation of stellar clusters as well as galaxies. We assume that the clouds condense out of a diffuse hot medium as a result of thermal instability. Large condensations with cooling timescales $\\tau_c > \\tau_d$ are thermally stable because they can adjust through contraction such that their radiative losses may be compensated by the release of their gravitational energy. Runaway cooling of the gas through thermal instability however occurs in clouds with $\\tau_c < \\tau_d$. In order to form clumps within an initially almost homogeneous cloud, internal density fluctuations must grow rapidly on a timescale short compared to the mean dynamical timescale of the entire cloud. Small scale density fluctuations would begin to dominate if either the growth timescale or the limiting amplitude is a decreasing function of the perturbations' length scale.  One possible fragmentation mechanism is gravitational collapse. The reduction in the cloud's temperature reduces its Jeans' mass, leading to the onset of gravitational instability and collapse.  However, for a non rotating, cold, homogeneous gaseous region, gravitational instability alone cannot induce fragmentation because the growth rate is essentially independent of length scale such that the growth timescale for the density contrast is comparable to the dynamical timescale of the whole cloud (Hunter 1962).  This has also been shown by numerical collapse simulations of initially gravitationally unstable perturbed gas clouds (e.g. Burkert \\& Bodenheimer 1993, 1996, Burkert, Bate \\& Bodenheimer 1997).  If the initial density perturbations $\\delta_0$ are linear ($\\delta_0 < 1$), fragmentation is suppressed until the gas cloud has collapsed into either a disk or a dense filamentary substructure.  We propose that clumpyness in clouds arises naturally from their formation through a cooling instability which acts on timescales that can be much shorter than the dynamical timescale of the cloud. In a pioneering paper, Field (1965) derived a criterion for a cooling gas to be unstable to the growth of thermal condensations. He showed that thermal instability can lead to the rapid growth of density perturbations from infinitesimal $\\delta_0$ to nonlinear amplitudes on a cooling timescale $\\tau_c$ which for typical conditions in the interstellar medium is short compared to the dynamical timescale.  If $\\tau_c$ increases with decreasing density any small density difference would induce a temperature difference between the cooler perturbed region and the warmer background. Across the interface between the two-phase medium, differential cooling leads to a pressure gradient which induces a gas flow from the lower-density background towards the higher-density perturbed region.  The density enhancement in the cooler region further reduces its cooling timescale compared to that of the background where $\\tau_c$ increases. A more detailed investigation of the growth of condensations in cooling regions has been presented by Schwarz et al. (1972) who included also the effects of ionization and recombination and by Balbus (1986) who examined the effect of magnetic fields. The classical model of the interstellar medium where heating balances cooling was presented by Field et al. (1969).  A recent progress report on the theory of thermal instability is given by Balbus (1995).  Although thermal instability proceeds faster than the collapse of the cloud, its growth rate is determined by the local cooling rate. During the initial linear evolution, variations in the initial over density (or under temperature) might lead only to a weak dependence of the growth timescale on the wavelength. In this paper we show however that there exist two important transitions which are very sensitively determined by the wavelengths of perturbations. 1) The growth of a perturbation is limited by its transition from isobaric to isochoric cooling, when the cooling time scale is reduced below the sound crossing time scale across the wavelength of the perturbation. This transition occurs at a lower temperature, with correspondingly larger over density, for perturbations with smaller wavelengths. 2) For those perturbation which can become nonlinear before the isobaric to isochoric transition, advection overtakes the pressure gradient in promoting the compression and growth of the perturbed region at an accelerated rate.  The fluctuations which can first reach nonlinearity would dominate the growth of all perturbations with longer wavelengths and homogenize disturbances with smaller wavelengths.  Thus, they determine the characteristic size and mass of the cold dense clumps which would emerge from the cooling of an initially nearly homogeneous cloud. Thermal conduction could in general erase these fluctuations, suppressing the instability. Weak, tangled magnetic fields would however be efficient enough in reducing the conductive flux, allowing the medium to break up into cold clumps on the characteristic length scale.  We study the cooling and fragmentation of gas using simplified power-law cooling functions. Since we are primarily interested in supercritical clouds, we neglect the effect of magnetic fields.  Note that even a weak magnetic field could have an important destabilizing influence in thermal instability (Loewenstein 1990, Balbus 1995). In \\S2, we obtain approximate analytic solutions which describe the evolution of a linear density perturbation in the isobaric and nearly isochoric regime.  We show that the growth of over density in a thermally unstable fluctuation is limited by a transition from isobaric to isochoric evolution and that the limiting amplitude is a decreasing function of the length scale. We verify our analytic approximations with numerical, hydrodynamical calculations which are also used in \\S3 to study the transition into the non-linear regime.  In \\S4 we investigate the cooling of interacting perturbations and determine the critical length scale of clumps that emerge through thermal instability.  The importance of thermal conduction is investigated in \\S5. In \\S6 we discuss the affect of heating processes and the formation of a stable 2-phase medium. Finally, we summarize our results and discuss their implications in \\S7. ",
        "conclusions": "The discussions in this paper focussed on the emergence of small scale perturbations.  We have assumed the pre-existence of small initial perturbations which is a reasonable assumption for dynamically evolving systems like the interstellar medium in galaxies or in galactic clusters.  We have limited our analysis to the optically thin regime such that radiation transfer is solely due to optically thin local radiative processes.  This approximation is appropriate for the collapse of supercritical clouds where the effect of a magnetic field is dominated by thermal processes.  Such a situation may be particularly relevant for the formation of stellar clusters and first generation stars in galaxies. Provided that the density of the progenitor clouds is relatively small, the local cooling approximation is adequate.  We also neglected the interaction and merging of clumps.  These processes become important for the subsequent evolution and they will be considered in subsequent papers.  In the context of our approximations, we have shown that thermal instability can lead to the breakup of large clouds into cold, dense clumps with a characteristic length scale which is given by $\\lambda_{crit}$ in eq. ({\\ref{a30}) or by the smallest unstable wavelength that is not erased by thermal conduction, depending on whether the amplitude of the initial perturbation is an increasing or decreasing function of wavelength. For linear perturbations with overdensities $\\rho_a/\\rho_0 \\approx 0.01$ the critical wavelength lies in the regime of $10^{-3}$ pc to $10^{-1}$ pc, depending on the initial temperature. The emergence of small scale dense subcondensations is equivalent to fragmentation. As in a thermally unstable region the cooling timescale is shorter than the dynamical timescale, gravity has no time to play an important role during this fragmentation process.  $\\lambda_{crit}$ may be either smaller or bigger than the Jeans' length. In the latter case gravity becomes important eventually. In general however, thermally induced fragmentation of clouds with small initial density fluctuations proceeds the onset of gravitational instability of their individual clumps.  In our analyses, we adopted an idealized power-law cooling function.  In reality, the cooling efficiency would terminate when the main cooling agents reach their ground state or establish an equilibrium with some external heating source. The latter is necessary for the clouds to attain a two-phase medium.  Interaction between these two phases may determine the pressure, density and infall rate of the cloud complex as well as the dynamical evolution and size distribution of cloudlets and sub condensations. The analysis of this interaction will be presented elsewhere.  \\newpage"
    },
    "0002/astro-ph0002276_arXiv.txt": {
        "abstract": "A search for new $\\gamma$ Dor stars was undertaken using the Hipparcos periodic variable star catalogue and the Geneva photometric database, leading to a list of 40 new candidates. We started a monitoring of the candidates which suited the observational window with the CORALIE spectrograph at the Swiss Euler Telescope for establishing a robust list of new $\\gamma$ Dor stars and studying line profile variations. We here present our long-term program. ",
        "introduction": "The $\\gamma$ Dor stars have amplitude variations up to 0.1 mag in Johnson V and periods ranging from 0.4 to 3 days (Kaye, these proceedings). We searched in two databases for finding new members of this class of variable stars. The first one is the Hipparcos main mission photometric database. It contains a mean of 110 measurements for 118\\,204 stars brighter than 12.4 and is magnitude complete up to 7.3-7.9 depending on the galactic latitude $b$. As the sampling is ruled by the scanning law of the satellite, it is not affected by the aliasing around 1/day, which might be a problem for detecting $\\gamma$ Dor stars.  The second scanned database is the Geneva photometric catalogue (Burki \\& Kienzle, these proceedings), it counts 48\\,000 stars and 345\\,000 measurements in a seven colour system. The content of the Geneva catalogue is the reunion of more than 200 scientific programmes, including namely the Bright Star Catalogue south of $\\delta < +20$. ",
        "conclusions": ""
    },
    "0002/astro-ph0002040_arXiv.txt": {
        "abstract": "We have begun an observational program to search nearby stars for dust disks that are analogous to the disk of zodiacal dust that fills the interior of our solar system. We imaged six nearby main-sequence stars with the Keck telescope at 11.6 microns, correcting for atmosphere-induced wavefront aberrations and deconvolving the point spread function via classical speckle analysis.  We compare our data to a simple model of the zodiacal dust in our own system based on COBE/DIRBE observations (Kelsall et al. 1998) and place upper limits on the density of exo-zodiacal dust in these systems. ",
        "introduction": " ",
        "conclusions": "To interpret our observations we compared them to models of the IR emission from the solar zodiacal cloud. We constructed a model for exo-zodiacal emission based on the smooth component of the Kelsall et al. \\markcite{kels98}(1998) model of the solar system zodiacal cloud as seen by COBE/DIRBE, with emissivity $\\epsilon \\propto r^{-0.34}$ and a temperature  $T = 286 \\ {\\rm K} \\ r^{-0.467}L^{0.234}$, where $r$ is the distance from the star in AU, and $L$ is the luminosity of the star in terms of $L_\\odot$. For a dust cloud consisting entirely of a single kind of dust particle of a given size and albedo, the $L$ exponent in the expression for the temperature is simply $-1/2$ times the $r$ exponent \\markcite{back93}(Backman \\& Paresce 1993).  The physics of the innermost part of the solar zodiacal dust is complicated (see \\markcite{mann93} Mann \\& MacQueen 1993), but our results are not sensitive to the details, because the hottest dust is too close to the star for us to resolve. We assume that the dust sublimates at a temperature of 1500 K, and allow this assumption to define the inner radius of the disk.  We set the outer radius of the model to 3 AU, the heliocentric distance of the inner edge of our own main asteroid belt. Our conclusions are not sensitive to this assumption; decreasing the outer radius to 2 AU or increasing it to infinity makes a negligible difference in the visibility of the model, even for A stars.  The assumed surface density profile, however, does make a difference. A collisionless cloud of dust in approximately circular orbits spiraling into a star due to Poynting-Robertson drag that is steadily replenished at its outer edge attains an equilibrium surface density that is independent of radius \\markcite{wyat50}\\markcite{brig62} (Wyatt and Whipple 1950, Briggs 1962). Models that fit data from the Helios space probes \\markcite{lein81}(Leinert et al 1981), the fit by Kelsall et al. \\markcite{kels98}(1998) to the COBE/DIRBE measurements and Good's \\markcite{good97}(1997) revised fit to the IRAS data all have surface densities that go roughly as $r^{-0.4}$.  This distribution appears to continue all the way in to the solar corona \\markcite{macq95}(MacQueen \\& Greely 1995).  We find that in general, if we assume an $r^{-\\alpha}$ surface density profile, our upper limit for the 1 AU density of a given disk scales roughly as $10^{\\alpha/2}$; disks with more dust towards the outer edge of the 11.6 micron emitting region are easier to resolve.  Likewise, the assumed temperature profile strongly affects our upper limits.  Unfortunately, we know little about the temperature profile of the solar zodiacal cloud.  COBE/DIRBE and IRAS only probed the dust thermal emission near 1 AU, and Helios measured the solar system cloud in scattered light, which does not indicate the dust temperature. We found that a dust cloud model with the IRAS temperature profile ($T = 266 \\ {\\rm K} \\ r^{-0.359}L^{0.180}$) was much easier to resolve than the model based on DIRBE measurements that we present here, especially for G and K stars.  To compare the models with the observations, we synthesized high resolution images of the model disks at an inclination of 30 degrees.  We calculated the IR flux of the stars from the blackbody function, and obtained the parallaxes of the stars from the Hipparcos Catalog \\markcite{esa97}(ESA 1997).   We inferred stellar radii and effective temperatures for each star from the literature and checked them by comparing the blackbody fluxes to spectral energy distributions based on photometry from the SIMBAD database \\markcite{simbad} (Egret et al. 1991).  For Altair and Vega, we use the interferometrically measured angular diameters \\markcite{hanb74}(1974) (they are 2.98 +/-0.14 mas and 3.24 mas). Stellar fluxes typically disagree with fitted blackbody curves by $\\sim10\\%$ in the mid-infrared \\markcite{enge90}(Engelke 1990), but our method does not require precise photometry, and the blackbody numbers suffice for determining conservative upper limits. We computed the power spectra of the images, and normalized them just like the observed power spectra.  In figures 3 and 4, the azimuthally-averaged power spectra for our target stars are compared to the extrapolated COBE/DIRBE model at a range of model surface densities.  Disks with masses as high as $10^3$ times the mass of the solar disk will suffer collisional depletion in their inner regions, so they are unlikely to have the same structure as the solar disk.  By neglecting this effect we are being conservative in our mass limits.  The density of the densest model disk consistent with the data in each case is listed in table 1.  Altair  Our best upper limit is for Altair (spectral type A7, distance 5.1 pc); with 11 pairs of object and calibrator observations we were able to rule out a solar-type disk a few times $10^3$ as dense as our zodiacal cloud.  Such a disk would have been marginally detectable by IRAS as a photometric excess.  Vega  IRAS detected no infrared excess in Vega's spectral energy distribution at 12 microns, with an uncertainty of 0.8 Jy. This may be due to a central void in the disk interior to about 26 AU \\markcite{back93} (Backman \\& Paresce 1993).  \\markcite{auma84}Aumann et al (1984) suggested that Vega (A0, 7.8 pc) could have a hot grain component (500 K) with up to $10^{-3}$ of the grain area of the observed component and not violate this limit.  The apparent upward trend in the visibility data may be a symptom of resolved flux in the calibrator stars.  We have only 3 object/calibrator pairs for Vega, not enough to test this hypothesis. Our upper limit is a solar-type disk with approximately $3 \\times 10^3$ times the density of the solar disk.  This disk would have a $\\geq$500 K emitting area of $10^{24} \\ {\\rm cm}^2$, about $10^{-3}$ of the grain area of the observed component.  61 Cygni A and B  Though 61 Cygni is close to the galactic plane and surrounded by cool cirrus emission, \\markcite{back86}Backman, Gillett and Low (1986) identified an IRAS point source with this binary system and deduced a far-infrared excess not unlike Vega's.  The color temperature of the excess suggests the presence of dust at distances $> 15$ AU from either star. However these stars are dim (spectral types K5 and K7) and the region of the disk hot enough to emit strongly at 11.6 microns is close to the star and difficult to resolve; we could not detect a solar-type dust disk around either of these objects at any density, assuming the COBE/DIRBE model, or unless it had $10^{5}$ times the density of the solar disk, assuming the IRAS model.  70 Oph B  70 Oph is a binary (types K0 and K4) with a separation of 24 pixels (2.6 arcsec).   We were able to assemble a power spectrum for B from 9 object/calibrator pairs, but the image of A fell on a part of the LWS chip that suffered from many bad pixels and was unusable.  The image of A may also have been distorted by off-axis effects.  70 Oph B, like 61 Cygni A and B, is dim, making any dust around it cool and hard to detect at 11.6 microns.  $\\tau$ Ceti  IRAS could have barely detected a disk with $\\sim 1000$ times the emitting area of the solar disk around Tau Ceti (G8, 3.6 pc), the nearest G star. We have only three object/calibrator pairs for this object, not enough data to improve on this limit."
    },
    "0002/astro-ph0002089_arXiv.txt": {
        "abstract": "The origin of the highest-energy cosmic rays remains elusive. The decay of a superheavy particle (X) into an ultra-energetic neutrino which scatters from a relic (anti-)neutrino at the Z-resonance has attractive features. Given the necessary X mass of $10^{14\\sim15}$ GeV, the required lifetime, $10^{15\\sim16}$ y, renders model-building a serious challenge but three logical possibilities are considered: (i) X is a Higgs scalar in $SU(15)$ belonging to high-rank representation, leading to {\\it power}-enhanced lifetime; (ii) a global X quantum number has {\\it exponentially}-suppressed symmetry-breaking by instantons; and (iii) with additional space dimension(s) localisation of X within the real-world brane leads to {\\it gaussian} decay suppression, the most efficient of the suppression mechanisms considered. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/hep-ph0002004_arXiv.txt": {
        "abstract": "It has been recently shown that it is possible to excite gravitinos in an expanding background due to time-varying scalar field oscillating at the bottom of the inflationary potential. The two components of the gravitino, namely helicity $1/2$ and helicity $3/2$ are excited differently due to the presence of different time-varying mass scales in the problem. In this paper we analyse the production of both the helicities in multi-chiral scenario, in particular focusing on a general model of hybrid inflation. Fermion production in hybrid models is very much different from that of the chaotic models discussed so far in the literature. In this paper we give a full account of gravitino production analytically and numerically.  It is noticed that the creation of gravitinos does not take place in the first few oscillations of the inflaton field, rather the production is a gradual and delayed process. It takes roughly $30-40$ oscillations to build up the production and for the saturation to take place it can even take longer time, depending on the model parameters.  We give an estimation of the reheat temperature and a brief discussion upon back-reaction on the gravitino production which could change its abundance. ",
        "introduction": "Low energy effective $N=1$ supergravity is a predictive theory \\cite{nilles}, which could predict an inflationary potential flat enough to provide adequate density perturbations \\cite{lyth}. So far, such viable inflationary models were constrained from observations by fixing the height of the potential, which essentially determines the amplitude of the COBE normalization. The first and the second derivative of the potential determines the tilt in the power spectrum, and the Yukawa couplings of the inflaton to other particles determines the reheat temperature of the Universe. The higher the coupling constant is, the higher is the temperature of the thermal bath and so the creation of the gravitinos from the collisions or the decay of other particles. The gravitinos decay very late and depending on their mass their life time could be long enough to disrupt the synthesis of light elements, via hadronic shower or by altering the entropy density of the baryons. Due to these reasons there is a strong constraint on the reheat temperature, for a review see Ref.~\\cite{sarkar}.   However, there is also a non-thermal phase of the Universe, just after the end of slow-roll inflation when the scalar field begins to oscillate coherently at the bottom of the potential. During this era an explosive production of particles, both bosons \\cite {linde} and fermions \\cite{patrick,peloso}, may take place due to the non-perturbative decay of the inflaton to other fields. It has also been shown that it is possible to create super-massive bosons and fermions. However, fermionic production is always saturated by the Pauli blocking. As a matter of fact creation of heavy non-thermal bosons can be a good candidate for weakly interacting massive particles, known as WIMPS \\cite{rocky}. The decay of super-massive bosons can explain the ultra-high energy cosmic rays \\cite{berezin}. Super-massive fermions can be used in leptogenesis, mainly from the decay of right-handed neutrinos into Higgs and leptons, which explicitly violates the CP conserving phase \\cite{peloso}. It is worth mentioning that the non-perturbative technique of decaying inflaton to other particles has given a new paradigm shift in understanding the hot big bang universe from the ultra-cold inflationary regime.   Recently preheating in the context of global supersymmetric theories has been considered \\cite{mar,mar1} and, as a natural extension it has been necessary to consider a local version of the supersymmetric theory and discuss the non-perturbative aspects of particle production and their consequences to nucleosynthesis. The local version of supersymmetry, known as supergravity, naturally accommodates the graviton and its superpartner the gravitino, a spin $3/2$ particle. Quantization of spin $3/2$ particles in the presence of an external background is plagued by consistency problems.  It has been known for a long time that quantization of spin $3/2$ particles in scalar, electromagnetic, or gravitational backgrounds can give rise to acausal behaviour \\cite{velo}. Supergravity is the only set-up where such problems do not occur, provided the background fields also satisfy the corresponding equations of motion \\cite{deser0}. Nevertheless, the complicated form of the Rarita-Schwinger equation makes it extremely difficult to extract any explicit results even in a simple background \\cite{auvil}. The problem was first addressed in Ref.~\\cite{deser}, where the authors quantized spin $3/2$ particles in a non-vanishing cosmological background, almost two decades ago.  The slightest generalization of quantizing spin $3/2$ has been done very recently in literature, in Ref.~\\cite{anupam}. The authors have extended the calculation of quantizing spin $3/2$ in presence of a time-varying homogeneously oscillating scalar field in a cosmological background.  This is the first result where the non-perturbative decay of the inflaton to gravitinos during preheating has been taken into account.  The authors have explicitly shown the production of a particular helicity, the $\\pm 3/2$ components of gravitino, in a particular new-inflationary type model \\cite{ross}. It has been noticed that the non-perturbative result can give rise to a larger abundance compared to the perturbative decay. The gravitino to photon number density has been found to be $n_{3/2}/n_\\gamma \\sim 10^{-12}$ \\cite{anupam}.  This abundance is $3$ orders of magnitude larger than the thermal abundance for a reheat temperature $10^{5}$ GeV \\cite{sarkar}. Such over production of gravitinos has been the first proof of the non-thermal production of gravitinos with helicity $\\pm 3/2$, which demands constraining the reheat temperature in any supergravity motivated inflationary model.   However, a massive gravitino has $4$ degrees of freedom, the other two degrees of freedom due to the helicity $\\pm 1/2$ components of the gravitino. The production of helicity $\\pm 1/2$ gravitinos is directly related to the problem of super-Higgs mechanism in supergravity models, which was studied in Ref.~\\cite{deser} in the context of a non-vanishing cosmological constant, and in Ref.~\\cite{cremmer} with a vanishing cosmological constant. In the presence of a time-varying scalar field there is an additional source of supersymmetry breaking via the non-vanishing time derivative of the homogeneous scalar field. This plays an important role in the context of cosmology when the scalar field is recognized as the inflaton, oscillating coherently at the bottom of the scalar potential. Due to the presence of such a field, supersymmetry is always broken at the minima of the potential and the initially massless gravitino which possesses only the helicity $\\pm 3/2$ components ``eats'' the goldstino to gain the other $\\pm 1/2$ components. At this point, one may wonder how to generalize the super-Higgs mechanism in such a scenario. In fact the problem turns out to be quite complicated and it has been addressed in two seminal papers \\cite{kallosh,riotto1}. These papers also study for the first time the production mechanism of the helicity $\\pm 1/2$ components of gravitino ( see other papers in the similar context \\cite{lyth1}). In all these papers the authors have assumed the existence of a unitary gauge, where the physical Lagrangian is free from the goldstino field. We will explicitly show that it is possible to choose such a gauge, where the gravitino equation of motion is free from the goldstino field. Our calculation is valid for more than one chiral field as well. Though we shall give a proof for the $F$-type supersymmetry breaking, this can be extended to $D$-type supersymmetry breaking also.   It has been noticed in Refs.\\cite{kallosh,riotto1} that the production of the two helicity states are completely different. The helicity $\\pm 1/2$ components are produced copiously compared to the helicity $\\pm 3/2$. In the case of helicity $\\pm 3/2$, conformal invariance is broken due to the presence of a time-varying gravitino mass, which is usually Planck mass suppressed, whereas for helicity $\\pm 1/2$  the breaking of conformal invariance is related to the presence of a massive Goldstone fermion \\cite{kallosh,riotto1}, whose time-varying mass is not suppressed by the Planck mass. Moreover in the high momentum limit the helicity $\\pm 1/2$ gravitino behaves like a fermion, the goldstino, as it is stated by the  equivalence principle \\cite{riotto2,maroto}. This has been studied in a single-chiral field scenario, where the source of conformal breaking can be directly related to the mass of Goldstino. Unfortunately in the multi-chiral field scenario the quantization scheme becomes more involved, and the relation between conformal breaking and Goldstino mass is not so straightforward. The situation has been briefly discussed in Ref. \\cite{riotto2}, and an attempt of a perturbative scheme has been suggested. However, it would be nice to discuss a non-perturbative scheme. Our paper fills that gap and, as we shall see, we can  discuss non-perturbative production of the helicity $\\pm 1/2$ gravitinos in the multi-chiral case.    The best example to study the multi-chiral field scenario is in the context of a general class of supersymmetric hybrid inflation model \\cite{mar,mar1}. There are essentially two scalar fields, one is responsible for inflation, and the other field is responsible for the phase transition which results in terminating the inflationary era. Unlike the non-supersymmetric version of hybrid inflation model \\cite{hybrid}, the supersymmetric version considered in this paper has only one coupling constant in the potential. This leads to a single natural frequency of oscillations. This gives us an ample opportunity to use techniques to explore gravitino production in a similar spirit as in the case of a single-chiral field. Fermionic production in this case is very much different compared to that of a quadratic inflationary potential. In the hybrid models the effective mass term for the fermions is always positive and as a result the production can never be completed in just a few oscillations, rather the occupation number gradually increases and depending on the model parameters the rate of production could be slowed down significantly. For example, in order to reach the Fermi saturation it may be necessary  more than $100$ oscillations. Due to such a slow rate of production, the issue of back-reaction becomes important. It is very likely that simultaneous non-perturbative production of bosons can change the picture quite significantly. In some sense, hybrid model can be considered to be the safest of all the supergravity motivated inflationary models, because the gravitino production can stop due to back-reaction effects coming from other newly created bosons or fermions. In other models, such as in the chaotic models, where most of the particle creation  takes place in the very first few oscillations, the issue of back-reaction hardly plays any significant role.  The layout of the paper is as follows; in section $2$, we establish the supergravity Lagrangian and discuss the gauge fixing mechanism for the multi-chiral field scenario; this is a mere generalization of the $R_{\\zeta}$ gauge usually used in pure gauge theories as well as an alternative way of explaining the existence of a unitary gauge in supergravity theories \\cite{baulieu,casalbuoni,maroto}.  In section $3$, we discuss the quantization procedure for the gravitinos. In section $4$, we briefly describe a general class of supersymmetric hybrid inflationary models we use. Analytical results and discussions based on our numerical results are presented in section $5$. The implications of these results on the reheat temperature are discussed in section $6$.  We give a detailed discussion on fermionic creation in hybrid model in the appendix. ",
        "conclusions": "We have carried out the calculation for the gravitino production in multi-chiral field scenario, in particular in the context of hybrid model. As we have shown, it is possible to add a gauge-fixing term to the supergravity Lagrangian to get rid of the mixing between the goldstino and the gravitino field. Our method is a generalization of $R_{\\zeta}$ gauge studied in various contexts. We choose to work in a unitary gauge, in which the goldstino is completely removed from the physical spectrum. Our study emphasizes major points in the non-perturbative production mechanism of gravitinos, analytically and numerically for the multi-chiral field models. We also give detailed analysis of the fermionic creation in general. We have observed that the fermionic creation in hybrid model is quite different from other chaotic inflationary models. The effective fermionic mass in hybrid models  never vanishes, and as a result the particle production does not take place in first few oscillations, rather it builds up gradually. This makes it more interesting as far as the gravitino production is concerned. If we really want nucleosynthesis to be preserved in the context of supergravity inflationary models, we believe models based on hybrid inflation with low scales are probably going to be the only saviour. The reason is very simple; in other models there is no way we can argue the back-reaction due to the creation of other particles would stop creating gravitinos, but in hybrid models there is a scope where the back-reaction due to non-perturbative creation of bosons could affect the coherent oscillations of the inflaton and halt the particle production completely. This gives us a new hope to understand the abundance of gravitinos during nucleosynthesis and we leave these important issues to be investigated in near future."
    },
    "0002/astro-ph0002531_arXiv.txt": {
        "abstract": "We present plane-parallel equilibrium models of molecular clumps that are supported by Alfv\\'en waves damped by the linear process of ion-neutral friction. We used a WKB approximation to treat the inward propagation of waves and adopted a realistic ionization structure influenced by dissociation and ionization due to photons of external origin. The model clumps are larger and less centrally condensed than those obtained for an assumed ionization structure, used in some previous studies, that is more appropriate for dark regions. ",
        "introduction": "Giant molecular cloud complexes (GMCs), the birth places of stars, are typically many tens of parsecs in linear extent and have masses from $10^4$ to $10^6 M_{\\sun}$ and temperatures of 10--30 K (see \\cite{HCRRW98} for a recent review). Observations of CO emission from GMCs (\\cite{BT80}; \\cite{WBS95}) show them to be composed of many smaller clumps that are a few parsecs in extent and contain $\\simlt 10^3 M_{\\sun}$.  The widths of CO emission lines originating in individual clumps are supersonic and have been attributed to the presence of Alfv\\'en waves having subAlfv\\'enic velocity amplitudes (\\cite{AM75}).  The Alfv\\'en waves contribute to the support of a clump along the direction of the large-scale magnetic field; the damping of the waves affects the degree of support that they provide.  An important and well understood mechanism for the damping of linear Alfv\\'en waves in a partially ionized medium is that due to ion-neutral friction which depends on the ionization structure (\\cite{KP69}).  Ruffle et al. (1998, hereafter R98) and Hartquist et al. (1993) have emphasised that the dependence of the ionization structure on total visual extinction, $A_\\mathrm{V}$, should greatly influence the density profiles of clumps if Alfv\\'en waves contribute to their support.  To quantify the assertion of Hartquist et al. (1993) and R98, we present in this paper models of plane-parallel, wave-supported GMC clumps like those identified in the work of Williams et al. (1995), who made a detailed analysis of the CO maps of the Rosette Molecular Cloud (RMC), identifying more than 70 clumps.  The models that we have constructed are for RMC-type clumps in equilibrium, a restriction justified by the fact that clear spectral signatures of collapse have been found only when much smaller scale features have been resolved (see, e.g., \\cite{HCRRW98}).  We have adopted a WKB description of the wave propagation as did Martin et al. (1997) in their work on wave-supported clumps.  Their work differs substantially from ours in that they used an ionization structure appropriate for dark regions.  Also, we have considered inwardly rather than outwardly propagating waves, as many of the clumps mapped by Williams et al. (1995) do not contain detected stars and may have no internal means of generating waves.  Indeed, the waves may be produced at the surface of a clump by its interaction with an interclump medium.  Other authors have addressed the importance of photoabsorption for the effects that the ionization structure will have upon a clump's dynamics.  These authors have been concerned primarily with dense cores and/or envelopes around them; cores are much smaller-scale objects than the clumps identified in Williams et al. (1995). McKee (1989) addressed the possibility that collapse in a system of dense cores is a self-regulating process due to the ionization of metals such as Magnesium and Sodium by photons emitted by stars formed in the collapse; he was concerned with infall due to ambipolar diffusion of a large-scale magnetic field.  Ciolek \\& Mouschovias (1995) have shown that the large-scale magnetic field can support a photoionized envelope around a dense core for a time that is very long compared to the ambipolar diffusion timescale in the center of the dense core.  In contrast to McKee (1989) and Ciolek \\& Mouschovias (1995), Myers \\& Lazarian (1998) addressed the effect of photoabsorption on support by waves rather than by the large-scale magnetic field.  They stressed that observed infall of dense core envelopes is slower than that expected due to gravitational free-fall and more rapid than collapse due to the reduction by ambipolar diffusion of support by an ordered large-scale magnetic field. They considered collapse of material supported primarily by waves and subjected to an external radiation field.  While they made clear comments about the importance of the $A_\\mathrm{V}$ dependence of the ionization structure for their model, they did not perform any detailed calculations in which a realistic dependence of the ionization fraction on $A_\\mathrm{V}$ was used.  Several sets of authors have considered nonlinear effects in the dissipation of waves supporting a clump.  Gammie \\& Ostriker (1996) investigated models of plane-parallel clumps and from their ``1 2/3-dimensional'' models found dissipation times due to nonlinear effects to be longer than the Alfv\\'en crossing times for a fairly large range of parameters. The three dimensional investigations of Mac Low et al. (1998) and Stone et al. (1998) suggest the more restrictive condition that the angular frequency of the longest waves be no more than a few times $2\\pi/t_\\mathrm{A}$ (where $t_\\mathrm{A}$ is the Alfv\\'en crossing time) in order for the dissipation timescale due to nonlinear damping to be roughly the Alfv\\'en crossing time or more.  In most cases addressed in this paper we have restricted our attention to such angular frequencies so that we are justified to lowest order in focusing on only the damping due to ion-neutral friction.  It should be noted that the above three dimensional studies of nonlinear effects concerned homogeneous turbulence and did not include ion-neutral damping for a realistic ionization structure. If we are correct in supposing that the waves in clumps are driven externally, then the turbulence is not homogeneous and its nature depends on both the viscous scale set by ion-neutral damping and the exact boundary conditions.  The effects of nonlinear damping and multiple dimensions will be considered in subsequent work.  In Sect. 2 we present the equations for the wave energy, the static equilibrium clump structure, and the gravitational field.  In Sect. 3 we give a description of the calculations of the ionization structure for various values of the clump density and $A_\\mathrm{V}$ while Sect. 4 contains details of the models considered here.  Finally, in Sect. 5, we present conclusions. ",
        "conclusions": "\\begin{figure} \\epsfxsize=8.8cm \\epsfbox{fig1_rho.eps} \\caption[]{Plot of density (normalized to the density at the outer edge of the clump) versus $A_\\mathrm{V}$ for the 5 Models.  The solid curve is for Model 1, the dotted curve for Model 2, the dashed curve for Model 3, the dot-dashed curve for Model 4, and the dash-chain-dot curve for Model 5.} \\label{fig:rho} \\end{figure}  \\begin{figure} \\epsfxsize=8.8cm \\epsfbox{fig2_U.eps} \\caption[]{As Fig.~\\ref{fig:rho} but for the perturbing magnetic field, $\\delta B$.} \\label{fig:deltaB} \\end{figure}  \\begin{figure} \\epsfxsize=8.8cm \\epsfbox{fig3_xi.eps} \\caption[]{As Fig.~\\ref{fig:rho} but for the absolute ion mass fraction, $\\xi$.  The bumps seen on the curves for Models 1--3 are a result of the interpolation over $A_\\mathrm{V}$.} \\label{fig:xi} \\end{figure}  \\begin{figure} \\epsfxsize=8.8cm \\epsfbox{fig4_flux.eps} \\caption[]{As Fig.~\\ref{fig:rho} but for the magnetic energy flux, $F \\equiv k_\\mathrm{r} U \\omega / |k|^2$.} \\label{fig:flux} \\end{figure}  \\begin{figure} \\epsfxsize=8.8cm \\epsfbox{fig5_zed.eps} \\caption[]{As Fig.~\\ref{fig:rho} but for z, the absolute spatial extent of the plane-parallel cloud.} \\label{fig:size} \\end{figure}  In Fig.~\\ref{fig:rho} we present density as a function of visual extinction for each of the 5 Models. It is clear that the newer ionization profiles used in Models 1--3 result in less condensed, more extended clumps than the profile used in Model 4.  In fact, $n_\\mathrm{H}$ is roughly proportional to $1/z$ in Models 1--3 while $n_\\mathrm{H}$ goes roughly as $1/z^2$ in Model 4.  In Fig.~\\ref{fig:deltaB} we show plots of $\\delta B$ versus $A_\\mathrm{V}$ for Models 1--4.  Except for Model 4, which has $k_\\mathrm{i}/k_\\mathrm{r}$ approaching $0.5$ at the clump boundary so that Eq.~\\ref{eq:test} is not satisfactorily satisfied, the perturbing field obeys flux conservation near the surface of the clump.  In Model 1, in the central region of the clump dissipation is rapid enough that $\\delta B$ begins to decrease.  In order to compensate for the loss of support, the equilibrium solution requires a complementary increase in the density, as can be seen in Fig.~\\ref{fig:rho}.  On the other hand, in Model 4, the turbulence is dissipated much nearer to the cloud boundary, thus requiring a steeper overall density profile. The higher ionization fractions of the case B depletions result in very little dissipation even in the center of the clump for Model 2. Note that even though observations (\\cite{WBS95}) suggest that the temperature of RMC-type clumps is closer to 10 K rather than the 20 K used here, thermal support is insignificant except in the centre of Model 1, so the effect of a lower temperature on the models would be merely to enhance slightly any central condensations.  The ionization profiles used in the Models are shown in Fig.~\\ref{fig:xi}.  The ionization profiles described in Sec. 3 result in $\\xi$ for Models 1--3 being more than 50 times greater near the surface of the clump than in Model 4.  However, in the center of the clump the ionization fraction drops, resulting in more dissipation.  Again, this leads to clumps that are overall more diffuse but with small condensed cores.  Clearly, clumps with $A_\\mathrm{V} \\simgt 5$ will have distinct central condensations with $n/n_\\mathrm{b} \\simgt 100$ and central fractional ionizations of $\\simlt 5\\times10^{-7}$.  Though dense cores may be formed during the fairly rapid collapse (as envisaged by Fielder \\& Mouschovias 1993) of more extended objects (i.e. RMC-like clumps) that become unstable, even in our equilibrium models we find central cores having densities and fractional ionizations similar to those measured for dense cores and their envelopes (\\cite{WBS95}; \\cite{WBCMP98}; \\cite{BPWM99}).  Fig.~\\ref{fig:flux} shows the flux of magnetic energy through the clumps for Models 1--4.  Near the clump center, Model 1 is nearly thermally supported due to the dissipation of the turbulence. The higher ionization fraction for the case B depletions used in Model 2 results in less dissipation and thus more turbulent support for the clump.  Consequently, as can be seen in Fig.~\\ref{fig:rho}, Model 2 has no central condensation and is more extended.  Unfortunately, we can only speculate about how the depletions of Sulphur, metals, and some other species behave in RMC-like clumps (\\cite{RHCW99}).  Thus, cases A and B are merely representative; as can be clearly seen in the figures, the clump profiles are very sensitive to the choice of abundances and the subsequent fractional ionizations. In addition, compared to Model 1, the stronger ion-neutral coupling in Model 3 results in less dissipation and subsequently the clump has little central condensation, as expected.  Note however that for Model 3 the lower limit in Eq.~\\ref{eq:omega} is not adequately satisfied.  Fig.~\\ref{fig:flux} also shows the effect of external wave generation.  If the fractional ionization is too low, as in Model 4, dissipation occurs close to the surface of the clump. Conversely, if the fractional ionization is too large, as in Model 2, significant dissipation occurs only at the clump's very centre.  Both extremes produce density profiles which lack a central condensation.  Note that if our externally generated wave model is correct, one should not see turbulence within a condensed core if there is no turbulence in its surrounding envelope.  In Fig.~\\ref{fig:size} we present curves which map the visual extinction to the spatial extent of the clumps.  Clouds with larger extents which match the observed 2--3 pc size of RMC-type clumps (\\cite{WBS95}) cannot be reproduced within the constraints given in Sect. 4.  However, observations generally measure the largest linear extent of a clump.   Thus, since the waves only support the model clumps parallel to the large-scale field, it is not surprising that the model sizes given here are less than the observed sizes.  Similarly, the models require high boundary densities and magnetic field strengths in order for the Alfv\\'en speed and wave velocity amplitude at visual extinctions where CO is abundant to be large enough to be compatible with observed linewidths. For Model 1, $n_\\mathrm{b}  = 375~\\H2$ cm$^{-3}$.  This is rather higher than the typical value of $n(\\H2) = 220$ cm$^{-3}$ given by Williams et al. (1995) for RMC-type clumps but, given the uncertainties, it is within a reasonable range of the Williams et al. (1995) value.  For Model 1, $B_0 = 135~\\mu$G, significantly higher than the value of $30\\, \\mu$G suggested by observations (\\cite{H87}) and expected from robust theoretical arguments (\\cite{M87}).  In order to determine whether the values of $n_\\mathrm{b}$ and $B_0$ could be lower and still allow model properties to be consistent with observed linewidths, we constructed models for clumps with total edge-to-center extinctions of 3 magnitudes.  The model giving $V = 2$ km sec$^{-1}$ and $v_\\mathrm{A} = 3$ km sec$^{-1}$ at $A_\\mathrm{V} = 2$ had $n_\\mathrm{b} = 325~\\H2$ cm$^{-3}$, $B_0 = 105\\,\\mu$G, $f_\\mathrm{b} = 0.49$, and $z_\\mathrm{max} = 0.565$ pc; although $n_\\mathrm{b}$ and $B_0$ were smaller and $z_{max}$ larger, the agreement with observations is nonetheless poor.  Thus, the next step in the modelling of clumps in which wave support is important is the inclusion of wave support in models analogous to the axisymmetric models of magnetically and thermally suported clumps described in classic papers by Mouschovias (1976a,1976b).  It is possible that the inclusion of magnetic tension, as well as pressure, will allow the reduction of $B_0$ to a value more like that expected and the construction of models of clumps having larger linear extents."
    },
    "0002/astro-ph0002477_arXiv.txt": {
        "abstract": "The normalized inclination distributions are presented for the spiral galaxies in RC3. The results show that, except for the bin of $81^{\\circ}$-$90^{\\circ}$, in which the apparent minor isophotal diameters that are used to obtain the inclinations, are affected by the central bulges, the distributions for Sa, Sab, Scd and Sd are well consistent with the Monte-Carlo simulation of random inclinations within 3-$\\sigma$, and Sb and Sbc almost, but Sc is different. One reason for the difference between the real distribution and the Monte-Carlo simulation of Sc may be that some quite inclined spirals, the arms of which are inherently loosely wound on the galactic plane and should be classified to Sc galaxies, have been incorrectly classified to the earlier ones, because the tightness of spiral arms which is one of the criteria of the Hubble classification in RC3 is different between on the galactic plane and on the tangent plane of the celestial sphere. Our result also implies that there might exist biases in the luminosity functions of individual Hubble types if spiral galaxies are only classified visually. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002194_arXiv.txt": {
        "abstract": "We consider resonant inverse Compton scattering of thermal photons by secondary particles above the pulsar polar gap. At neutron star temperatures higher than $10^5$ K the process appears to be an essential energy loss mechanism for the particles. The distribution function of the secondary plasma particles is found to be strongly affected by the scattering. It becomes two-humped implying the development of two-stream instability. The resonantly upscattered Compton photons are found to gain energy of 1--10 MeV forming an additional component in the pulsar gamma-ray spectrum. The corresponding gamma-ray flux is estimated as well. ",
        "introduction": "Rotation of a highly magnetized neutron star is known to induce a strong electric field, which intensely accelerates charged particles. According to the customary polar gap models (Ruderman \\& Sutherland 1975, Arons \\& Scharlemann 1979), the acceleration takes place near the neutron star surface above the polar cap, the Lorentz-factor of the primary particles increasing up to $\\sim 10^6$. The particles move along the magnetic lines and emit curvature photons, which initiate pair-production cascade. The first electron-positron pairs created screen the accelerating electric field, so that at higher altitudes the particle energy remains unaltered; the typical Lorentz-factors of the secondary plasma are $\\sim 10-10^4$.  Recent observations testify to thermal soft X-ray emission from some pulsars indicating that the neutron stars can be rather hot, $T\\sim 5\\cdot 10^5$ K, while the polar caps can have still higher temperatures scaling a few times $10^6$ K (Cordova {\\it et al.} 1989, Ogelman 1991, Halpern \\& Holt 1992, Finley {\\it et al.} 1992, Halpern \\& Ruderman 1993, Ogelman \\& Finley 1993, Ogelman {\\it et al.} 1993, Yancopoulos {\\it et al.} 1994, Ogelman 1995, Greiveldinger {\\it et al.} 1996). Such high temperatures of the neutron star surface are also predicted theoretically (Alpar {\\it et al.} 1984, Shibazaki \\& Lamb 1989, Van Riper 1991, Page \\& Applegate 1992, Umeda {\\it et al.} 1993, Halpern \\& Ruderman 1993). The thermal X-ray photons should suffer inverse Compton scattering off the primary particles in the polar gap. In the presence of a strong magnetic field the scattering cross-section is essentially enhanced, if the photon energy in the particle rest frame equals the cyclotron energy (Herold 1979, Xia {\\it et al.} 1985). For the pulsars with hot polar caps the resonant Compton scattering in the polar gap was found to be rather efficient (Kardashev {\\it et al.} 1984, Xia {\\it et al.} 1985, Daugherty \\& Harding 1989, Dermer 1990, Sturner 1995, Chang 1995). Firstly, it was recognized as an essential mechanism for energy loss of primary particles accelerating in the polar gap. Secondly, inverse Compton scattering was found to condition the gap formation, since the pair production avalanche may be triggered by the upscattered Compton photons rather than by curvature photons (Zhang \\& Qiao 1996, Qiao \\& Zhang 1996, Luo 1996, Zhang {\\it et al.} 1997).  As shown by Sturner (1995), given typical values of neutron star temperature and surface magnetic field, the resonant Compton scattering is the strongest for particle Lorentz-factors $\\sim 10^2-10^3$. Theoretical models (Van Riper 1991, Page \\& Applegate 1992, Umeda {\\it et al.} 1993) suggest that the typical temperatures of the neutron star surface are as high as a few times $10^5$ K. Hence, the scattering is likely to be essential for the secondary plasma in most of the pulsars. Daugherty \\& Harding (1989), Zhang {\\it et al.} (1997) traced the evolution of the Lorentz-factor of secondary particles on account of magnetic inverse Compton scattering above the polar gap. Our aim is to investigate in more detail the influence of resonant inverse Compton scattering on the parameters of secondary plasma. For the resonant character of the scattering, the evolution of particle Lorentz-factor with the distance depends strongly on the initial particle energy. Since the distribution function of the secondary plasma is generally believed to be rather broad (the Lorentz-factor ranges from $10$ to $10^4$), its evolution on account of the Compton scattering is of a great interest. It will be shown that only particles with Lorentz-factors between 100 and 1000 are essentially decelerated in the course of the resonant scattering forming a sharp peak at low energies. Particles with larger Lorentz-factors are not decelerated at all. Thus, the resultant distribution function of secondary particles becomes two-humped, giving rise to the two-stream instability.  In Sect. 2 we examine how the resonant inverse Compton scattering affects the Lorentz-factor of a secondary particle at various pulsar parameters. Section 3 is devoted to studying the evolution of the distribution function. The conditions for the development of the two-stream instability are also discussed. In Sect. 4 we estimate the gamma-ray luminosity caused by the upscattered Compton photons. Section 5 contains a brief summary. ",
        "conclusions": "We have investigated resonant inverse Compton scattering by secondary pulsar plasma. The process is found to cause the efficient energy loss of the secondary particles given the neutron star temperatures $>10^5$ K, so that our results are applicable to most pulsars. For the resonant character of the scattering, the energy loss depends strongly on the initial particle energy. At $\\gamma_0\\sim 10^2-10^3$ the scattering is the most essential, the final Lorentz-factors of the particles being independent of the initial ones.  The distribution function of the secondary plasma is significantly altered by the resonant inverse Compton scattering. It is shown that ultimately the distribution function becomes two-humped. The main peak at $\\gamma\\sim 10^2$ is very sharp. It is formed by particles which suffered severe energy loss on account of the scattering. Another hump is sufficiently broad. It is associated with the particles, whose Lorentz-factors are almost unaltered by the scattering. The two-humped distribution function of the plasma particles is known to be unstable. It is shown that at pulsar conditions the two-stream instability develops readily and leads to an essential increase of plasma oscillations, which are likely to be transformed into radio emission.  We have also estimated the gamma-ray flux provided by the upscattered Compton photons. The resonantly scattered photons appear to gain the energies of 1--10 MeV forming an additional low-energy component in pulsar gamma-ray spectrum."
    },
    "0002/astro-ph0002127_arXiv.txt": {
        "abstract": "The late-time entropy production by the massive particle decay induces the various cosmological effects in the early epoch and modify the standard scenario. We investigate the thermalization process of the neutrinos after the entropy production by solving the Boltzmann equations numerically.  We find that if the large entropy are produced at $t \\sim 1$ sec, the neutrinos are not thermalized very well and do not have the perfect Fermi-Dirac distribution. Then the freeze-out value of the neutron to proton ratio is altered considerably and the produced light elements, especially $\\4he$, are drastically changed. Comparing with the observational light element abundances, we find that $T_R \\lesssim 0.7$~MeV is excluded at 95 $\\%$ C.L. We also study the case in which the massive particle has a decay mode into hadrons. Then we find that $T_R$ should be a little higher, {\\it i.e.} $T_R \\gtrsim$ 2.5 MeV - 4 MeV, for the hadronic branching ratio $B_h = 10^{-2}-1$. Possible influence of late-time entropy production on the large scale structure formation and temperature anisotropies of cosmic microwave background is studied. It is expected that the future satellite experiments (MAP and PLANCK) to measure anisotropies of cosmic microwave background radiation temperature can detect the vestige of the late-time entropy production as a modification of the effective number of the neutrino species $N_{\\nu}^{\\rm eff}$. ",
        "introduction": "In the standard big bang cosmology it had been assumed tacitly that the universe was dominated by the thermal radiation at the early epoch.  Even in the paradigm of the modern cosmology it is commonly believed that thermal radiation was produced by the reheating process after the primordial inflation and they dominated the energy of the universe at sufficiently early epoch.  Here we ask, ``How early should the universe be dominated by radiation in order to success the standard big bang cosmology?''. We could say that the energy of the universe should be dominated by the radiation at least before the beginning of the big bang nucleosynthesis (BBN) epoch.  In this paper we answer the above question.  The various models of the modern particle physics beyond the standard model predicts a number of unstable massive particles which have long lifetimes and decays at about BBN epoch. The energy density of the non-relativistic particles or the oscillation energy density of the scalar fields (inflaton and so on) decreases as $\\rho_{NR}(t) \\propto a(t)^{-3}$, where $a(t)$ is a scale factor. On the other hand since the radiation energy density decreases more rapidly $\\rho(t)\\propto a(t)^{-4}$, if the energy density of the massive non-relativistic particles or the oscillating scalar fields is large enough, it immediately dominates the universe as it expands, and the universe necessarily becomes matter-dominated until the cosmic time reaches to their lifetime. When the particles decay into standard particles (e.g. photon and electron), they produce the large entropy and the universe becomes radiation-dominated again. It is expected that such process would change the initial condition for the standard big bang scenario. We call the process ``late-time entropy production''.  Now we have some interesting candidates for late-time entropy production in models based on supersymmetry (SUSY).  It is known that gravitino and Polonyi field which exist in local SUSY ({\\it i.e.} supergravity ) theories have masses of $\\sim {\\cal O}(100{\\rm GeV}-10{\\rm TeV})$~\\cite{Nilles}. In addition they have long lifetimes because they interact with the other particle only through gravity. For example, since Polonyi field~\\cite{Polonyi} which has a heavy mass $\\sim 10$~TeV cannot be diluted by the usual inflation, it immediately dominates the universe and decays at the BBN epoch. Moreover it is also known that in the superstring theories there exist many light fields called dilaton and moduli which have similar properties to the Polonyi field.  Recently Lyth and Stewart~\\cite{Lyth} considered a mini-inflation called ``thermal inflation'' which dilutes the above dangerous scalar fields. In the thermal inflation scenario, however, the flaton field which is responsible for the thermal inflation decays at late times. In particular, if Polonyi (moduli) mass is less than $\\sim 1$~GeV which is predicted in the framework of gauge-mediated SUSY breaking models~\\cite{Giudice}, the sufficient dilution requires that the flaton decays just before BBN~\\cite{Asaka}.  Thus, in thermal inflation models, one should take care of the late-time entropy production.  To keep the success of BBN, any long-lived massive particles or the coherent oscillation of any scalar fields which dominate the universe at that time must finally decay into the standard particles before the beginning of BBN. Moreover the decay products would have to be quickly thermalized through scatterings, annihilations, pair creations and further decays and make the thermal bath of photon, electron and neutrinos. Concerning photons and electrons which electromagnetically interact, the interaction rate is much more rapid than the Hubble expansion rate at that time. Therefore it is expected that the photon and electron which are produced in the decay and subsequent thermalization processes are efficiently thermalized. The problem is that neutrinos can interact only through the weak interaction. In the standard big bang cosmology the neutrinos usually decouple from the electromagnetic thermal bath at about $T \\simeq 2-3$MeV. Therefore it is approximately inferred that they can not be sufficiently thermalized at the temperature $T \\lesssim $ a few MeV. Namely the reheating temperature after the entropy production process should be high enough to thermalize the neutrinos. Though people had ever used the rough constraints on the reheating temperature between 1MeV - 10MeV, in the previous paper~\\cite{kks} we pointed out that the neutrino thermalization is the most crucial for the successful BBN. In this paper we describe the detail of the method to obtain the neutrino spectrum and the formulations to integrate a set of Boltzmann equations numerically , and we study the constraint on the reheating temperature using the obtained neutrino spectrum and the full BBN network calculations with the revised observational light element abundances.  The above constraint is almost model-independent and hence conservative because we only assume that the massive particle decay produces the entropy. However, a more stringent constraint can be obtained if we assume a decay mode into quarks or gluons. In this case some modifications are needed for the above description.  When the high energy quark-antiquark pairs or gluons are emitted, they immediately fragment into a lot of hadrons (pions , kaons, protons, neutrons, {\\it etc}.). It is expected that they significantly influence the freeze-out value of neutron to proton ratio at the beginning of BBN through the strong interaction with the ambient protons and neutrons.  In the previous paper~\\cite{kks} we did not consider such hadron injection effects on BBN. Therefore we carefully treat the hadron injection effects in the present paper.  For another constraint, the late-time entropy production may induce the significant effects on the anisotropies of the cosmic microwave background radiations (CMB). Lopez {\\it et al.}~\\cite{Lopez} pointed out that the CMB anisotropies are very sensitive to the equal time of matter and radiation. When the reheating temperature is so low that neutrinos do not be sufficiently thermalized, the radiation density which consists of photon and neutrinos becomes less than that in the standard big bang scenario. It may give distinguishable signals in the CMB anisotropies as a modification of the effective number of neutrino species $N_{\\nu}^{\\rm eff}$. With the present angular resolutions and sensitivities of COBE observation~\\cite{COBE} it is impossible to set a constraint on $N_{\\nu}^{\\rm eff}$ but it is expected that future satellite experiments such as MAP~\\cite{MAP} and PLANCK~\\cite{PLANCK} will gives us a useful information about $N_{\\nu}^{\\rm eff}$.  In addition the above effect may also induce the signals in the observed power spectrum of the density fluctuation for the large scale structure as a modification of the epoch of the matter-radiation equality.   The paper is organized as follows. In Sec.~II we introduce the formulation of the basic equations and the physical parameters. In Sec.~III we briefly review the current status of the observational light element abundances. In Sec.~IV we study the spectra of the electron neutrino and the mu(tau)-neutrino by numerically solving the Boltzmann equations, and the constraints from BBN are obtained there. In Sec~V we investigate the additional effects in the hadron injection by the massive particle decay.  In Sec.~VI we consider the another constraints which come from observations for large scale structures and anisotropies of CMB.  Sec~VII is devoted to conclusions. In Appendix we introduce the method of the reduction for the nine dimension integrals into one dimension. ",
        "conclusions": "\\label{sec:conclude} In this paper we have investigated the various cosmological effects induced by the late-time entropy production due to the massive particle decay. The neutrino distribution functions have been obtained by solving the Boltzmann equations numerically. We have found that if the large entropy are produced at about $t \\simeq 1$ sec, the neutrinos are not thermalized very well and hence do not have the perfect Fermi-Dirac distribution. The deficits of the neutrino distribution functions due to the insufficient thermalization decrease the Hubble expansion rate and weakens the weak interaction rates between proton and neutron.  The above two effects changes the freeze-out value of $n/p$ significantly.  Especially the produced $\\4he$ mass fraction $Y$ is so sensitive to $n/p$ that the predicted value of $Y$ is changed drastically.  Comparing the theoretical predictions of D, $\\4he$ and $\\li7$ to the observational data, we have estimated the lower bound on the reheating temperature $T_R$ after the entropy production.  We have found that $T_R \\lesssim 0.7$~MeV is excluded at 95 $\\%$ C.L. In other wards, $T_R$ can be as low as 0.7 MeV. Then the effective number of neutrino species $N_{\\nu}^{\\rm eff}$ can be as small as $0.1$. It is enough sensitive for the ongoing large scale structure observations such as 2DF and SDSS or future satellite experiments (MAP and PLANCK) of CMB anisotropies to detect such modifications on $N_{\\nu}^{\\rm eff}$ and we can find out the vestige of the late-time entropy production.  Furthermore, we have also studied the case in which the massive particle has some decay modes into quarks or gluons. In this scenario, a lot of hadrons, {\\it e.g.}  pions, kaons, protons and neutrons, which are originated by the fragmentation of the high energy quarks and gluons are injected into thermal bath. The emitted hadrons extraordinarily inter-convert the ambient protons and neutrons each other through the strong interaction even after the freeze-out time of the neutron to proton ratio $n/p$. Then the predicted value of $Y$ increase extremely and we can constrain $T_R$ and the branching ratio of the hadronic decay mode $B_h$ comparing to the observational light element abundances. We have found $T_R$ should be higher than 2.5 MeV - 4 MeV at 95 $\\%$ C.L. for $B_h$ = $10^{-2}$ - 1. The above results tells us that $N_{\\nu}^{\\rm eff}$ can be as small as 1.9 - 2.8 even in the hadron injection scenario for $B_h = 10^{-2}$ - 1. Then it still may be possible to detect the modifications on $N_{\\nu}^{\\rm eff}$ by MAP and PLANCK."
    },
    "0002/astro-ph0002257_arXiv.txt": {
        "abstract": "We implement an Independent Component Analysis (ICA) algorithm to separate signals of different origin in sky maps at several frequencies. Due to its self-organizing capability, it works without prior assumptions either on the frequency dependence or on the angular power spectrum of the various signals; rather, it learns directly from the input data how to identify the statistically independent components, on the assumption that all but, at most, one of them have non-Gaussian distributions.  We have applied the ICA algorithm to simulated patches of the sky at the four frequencies (30, 44, 70 and 100 GHz) of the Low Frequency Instrument (LFI) of ESA's {\\sc Planck} satellite. Simulations include the Cosmic Microwave Background (CMB), the synchrotron and thermal dust emissions and extragalactic radio sources. The effects of detectors angular response functions and of instrumental noise have been ignored in this first exploratory study. The ICA algorithm reconstructs the spatial distribution of each component with rms errors of about 1\\% for the CMB and of about $10\\%$ for the, much weaker, Galactic components. Radio sources are almost completely recovered down to a flux limit corresponding to $\\simeq 0.7\\sigma_{CMB}$, where $\\sigma_{CMB}$ is the rms level of CMB fluctuations. The signal recovered has equal quality on all scales larger then the pixel size. In addition, we show that for the strongest components (CMB and radio sources) the frequency scaling is recevered with percent precision. Thus, algorithms of the type presented here appear to be very promising tools for component separation. On the other hand, we have been dealing here with an highly idealized situation. Work to include instrumental noise, the effect of different resolving powers at different frequencies and a more complete and realistic characterization of astrophysical foregrounds is in progress. ",
        "introduction": "\\label{introduction}  Maps produced by large area surveys aimed at imaging primordial fluctuations of the Cosmic Microwave Background (CMB) contain a linear mixture of signals by several astrophysical and cosmological sources (Galactic synchrotron, free-free and dust emissions, both from compact and diffuse sources, extragalactic sources, Sunyaev-Zeldovich effect in clusters of galaxies or by inhomogeneous re-ionization, in addition to primary and secondary CMB anisotropies) convolved with the spatial and spectral responses of the antenna and of the detectors. In order to exploit the unique cosmological information encoded in the CMB anisotropy patterns as well as the extremely interesting astrophysical information carried by the foregound signals, we need to accurately separate the different components.  A great deal of work has been carried out in recent years in this area (see de Oliveira-Costa \\& Tegmark 1999, and references therein; Tegmark et al. 2000). The problem of map denoising has been tackled with the wavelets analysis on the whole sphere \\cite{TENORIO} and on sky patches \\cite{SANZb}. Algorithms to single out the CMB and the various foregrounds have been developed \\cite{WF,HOBSON,TE}. In these works, Wiener filtering (WF) and the maximum entropy method (MEM) have been applied to simulated data from the {\\sc Planck} satellite, taking into account the expected performances of the instruments. Assuming a perfect knowledge of the frequency dependence of all the components, as well as priors for the statistical properties of their spatial pattern, these algorithms are able to recover the the strongest components, at the best {\\sc Planck} resolution.  We adopt a rather different approach, considering denoising and deconvolution of the signals on one side and component separation on the other as separate steps in the data analysis process, and focus here on the latter step only, presenting a 'blind separation' method, based on 'Independent Component Analysis' (ICA) techniques. The method does not require any a priori assumption on spectral properties and on the spatial distribution of the various components, but only that they are statistically independent and all but at most one have a non-Gaussian distribution. It is important to note that this is in fact the physical system we have to deal with: surely all the foregrounds are non-Gaussian, while the CMB is expected to be a nearly Gaussian fluctuation field for most of the candidate theories of the early universe.  The paper is organized as follows. In Section 2 we introduce the relevant formalism and briefly review methods applied in previous works. In Section 3 we outline the ICA algorithm in a rather general framework, since it may be useful for a variety of astrophysical applications. In Section 4 we describe our simulated maps. In Section 5 we give some details on our analysis and present the results. In Section 6 we draw our conclusions and list some future developments. ",
        "conclusions": "We have developed a neural network suitable to implement the Independent Component Analysis technique for separating different emission components in maps of the sky at microwave wavelengths. The algorithm was applied to simulated maps of a $15^{\\circ}\\times 15^{\\circ}$ region of sky at 30, 44, 70, 100 GHz, corresponding to the frequency channels of {\\sc Planck}'s Low Frequency Instrument (LFI).  Simulations include the Cosmic Microwave Background, extragalactic radio sources and Galactic synchrotron and thermal dust emission. The various components have markedly different angular patterns, frequency dependences and amplitudes.  The technique exploits the statistical independence of the different signals to recover each individual component with no prior assumption either on their spatial pattern or on their frequency dependence. The great virtue of this approach is the capability of the algorithm to {\\it learn} how to recover the independent components in the input maps. The price of the lack of {\\it a priori} information is that each signal can be recovered multiplied by an unknown constant produced during the learning process itself. However this is not a substantial limitation, since a lot of physics is encoded in the spatial patterns of the signals, and ultimately the right normalization of each component can be obtained by resorting to independent observations.  The results are very promising. The CMB map is recovered with an accuracy at the 1\\% level. The algorithm is remarkably efficient also in the detection of extragalactic radio sources: almost all sources brighter then 15 mJy at 100 GHz (corresponding to $\\simeq 0.7 \\sigma_{CMB}$, $\\sigma_{CMB}$ being the rms level of CMB fluctuations on the pixel scale) are recovered; on the other hand, it must be stressed that is not directly indicative of what can be achieved in the analysis of Planck/LFI data because the adopted resolution ($3'.5\\times 3'.5$) is much better than that of the real experiment, instrumental noise has been neglected and the same spectral slope was assumed for all sources.  Also the frequency dependences of the strongest components are correctly recovered (error on the spectral index of 1\\% for the CMB and extragalactic sources).  Maps of subdominant signals (Galactic synchrotron and dust emissions) are recovered with rms errors of about 10\\%; their spectral properties cannot be retrieved by our technique.  The reconstruction has equal quality on all the scales of the input maps, down to the pixel size.  All this indicates that this technique is suitable for a variety of astrophysical applications, i.e. whenever we want to separate independent signals from different astrophysical processes occurring along the line of sight.  Of course, much work has to be done to better explore the potential of the ICA technique. It has to be tested under more realistic assumptions, taking into account instrumental noise and the effect of angular response functions as well as including a more complete and accurate characterization of foregrounds.  In particular, the assumption that the spectral properties of each foreground component is independent of position will have to be relaxed to allow for spectral variations across the sky. Also, it will be necessary to deal with the fact that Galactic emissions are correlated.  The technique is flexible enough to offer good prospects in this respect. In the learning stage, the ICA algorithm makes use of non-linear functions that, case by case, are chosen to minimize the mutual information between the outputs; improvements could be obtained by specializing the ICA inner non-linearities to our specific needs. Also, it is possible to take properly into account our prior knowledge on some of the signals to recover, still taking advantage as far as possible of the ability of this neural network approach to carry out a ``blind\" separation. Work in this direction is in progress.  \\vskip .1in We warmly thank Luigi Danese for original suggestions. We also thank Krzysztof M. G\\'orski and all the people who collaborated to build the HEALPix pixelization scheme extensively used in this work. Work supported in part by ASI and MURST. LT acknowledges financial support from the Spanish DGES, projects ESP98--1545--E and PB98--0531--C02--01."
    },
    "0002/astro-ph0002243_arXiv.txt": {
        "abstract": "The $\\gamma$-ray burst (GRB) model for production of ultra-high-energy, $>10^{19}$~eV, cosmic-rays is based on the hypothesis that GRBs arise from the dissipation of the kinetic energy of relativistic fireballs at cosmological distances. Recent observations of delayed low energy emission, ``afterglow,'' from GRB sources strongly support the validity of this hypothesis. Observations also provide quantitative support for the model. The inferred physical fireball parameters imply that protons may be accelerated to $>10^{20}$~eV, and the inferred GRB energy generation rate is similar to that required to account for the observed flux of ultra-high-energy cosmic-rays (UHECRs).  Strong suppression of cosmic-ray flux is expected in this model above $10^{19.7}$~eV, due to proton interaction with microwave background photons. Strong deviations from model flux derived under the assumption of uniform source distribution is expected above $10^{20}$~eV, due to source discreteness and due to inhomogeneities in source distribution. In particular, the flux above $10^{20.5}$~eV is expected to be dominated by few, narrow spectrum sources. While model predictions can not be tested (with high confidence level) using present data, the predicted signatures should be observed with the planned Auger and Telescope-Array UHECR detectors.  A natural consequence of the GRB model of UHECR production is the conversion of a large fraction, $\\sim10\\%$, of the fireball energy to accompanying burst of $\\sim10^{14}{\\rm eV}$ and $\\sim10^{18}{\\rm eV}$ neutrinos. A ${\\rm km}^2$ neutrino detector would observe several tens of events per year correlated with GRBs, and test for neutrino properties (e.g. flavor oscillations, for which upward moving $\\tau$'s would be a unique signature, and coupling to gravity) with an accuracy many orders of magnitude better than is currently possible. ",
        "introduction": "The origin of GRBs, bursts of 0.1 MeV---1 MeV photons lasting for a few seconds, remained unknown for over 20 years, primarily because GRBs were not detected prior to 1997 at wave-bands other than $\\gamma$-rays (see \\cite{Fishman} for review of $\\gamma$-ray observations). The isotropic distribution of bursts over the sky suggested that GRB sources lie at cosmological distances, and general phenomenological considerations were used to argue that the bursts are produced by the dissipation of the kinetic energy of a relativistic expanding fireball (see \\cite{fireballs} for review).  Adopting the cosmological fireball hypothesis, it was shown that the physical conditions in the fireball dissipation region allow Fermi acceleration of protons to energy $>10^{20}{\\rm eV}$ \\cite{W95a,Vietri95}, and that the average rate at which energy is emitted as $\\gamma$-rays by GRBs is comparable to the energy generation rate of UHECRs in a model where UHECRs are produced by a cosmological distribution of sources \\cite{W95b}. Based on these two facts, it was suggested that GRBs and UHECRs have a common origin (see \\cite{Nobel_rev} for review).  In the last two years, afterglows of GRBs have been discovered in X-ray, optical, and radio wave bands (see \\cite{AG_review} for review). Afterglow observations confirmed the cosmological origin of the bursts, through the redshift determination of several GRB host-galaxies (see \\cite{Freedman} for an updated list), and confirmed \\cite{AG_confirm} standard model predictions \\cite{AG_pred} of afterglows that result from the collision of an expanding fireball with its surrounding medium. These observations therefore provide strong support for the GRB model of UHECR production.  In this review, UHECR and neutrino production in GRBs is discussed in the light of recent GRB and UHECR observations. The fireball model is briefly described in \\S2.1, and proton acceleration in GRB fireballs is discussed in \\S2.2. Recent claims, according to which protons can not be accelerated to $>10^{20}$~eV in the fireball \\cite{Gallant98}, are shown in \\S2.2 to be erroneous. Implications of recent afterglow observations to high energy particle production are discussed in \\S3. It is shown that, contrary to some recent claims \\cite{Stecker}, the GRB energy generation rate implied by afterglow observations is similar to the energy generation rate required to account for the flux of $>10^{19}$~eV cosmic-rays. Model predictions are shown to be consistent with the observed UHECR spectrum in \\S4.  Predictions of the GRB model for UHECR production, that can be tested with future UHECR experiments, are discussed in \\S5. Implications of the detection by the AGASA experiment of multiple high energy events with consistent arrival directions \\cite{AGASA_pairs} is also discussed in \\S5. High energy neutrino production in fireballs and its implications for future high energy neutrino detectors are discussed in \\S6. ",
        "conclusions": ""
    },
    "0002/astro-ph0002305_arXiv.txt": {
        "abstract": "Recently correlation analyses between Galactic dust emission templates and a number of CMB data sets have led to differing claims on the origin of the Galactic contamination at low frequencies.  de Oliviera-Costa {\\em et al} (1999) have presented work based on Tenerife data supporting the spinning dust hypothesis.  Since the frequency coverage of these data is ideal to discriminate spectrally between spinning dust and free-free emission, we used the latest version of the Tenerife data, which have lower systematic uncertainty, to study the correlation in more detail.  We found however that the evidence in favor of spinning dust originates from a small region at low Galactic latitude where the significance of the correlation itself is low and is compromised by systematic effects in the Galactic plane signal.  The rest of the region was found to be uncorrelated.  Regions that correlate with higher significance tend to have a steeper spectrum, as is expected for free-free emission.  Averaging over all correlated regions yields dust-correlation coefficients of $180\\pm47$ and $123\\pm16$ $\\mu$K /MJy sr$^{-1}$ at 10 and 15~GHz respectively. These numbers however have large systematic uncertainties that we have identified and care should be taken when comparing with results from other experiments.  We do find evidence for synchrotron emission with spectral index steepening from radio to microwave frequencies, but we cannot make conclusive claims about the origin of the dust-correlated component based on the spectral index estimates.  Data with higher sensitivity are required to decide about the significance of the dust-correlation at high Galactic latitudes and other Galactic templates, in particular H$_\\alpha$ maps, will be necessary for constraining its origin. ",
        "introduction": "\\label{intro} A small, but significant correlation of existing Cosmic Microwave Background (CMB) data with maps of Galactic dust emission has been detected. It is important to understand the origin, and hence the characteristics of any Galactic emission present in CMB maps with structure on angular scales relevant to CMB measurements. With such information, we can attempt to remove these contaminating emissions from the data, to be left with a pure CMB map.  Cross-correlating the COBE DMR maps with DIRBE far-infrared maps, Kogut {\\em et al} (1996a,b) discovered that statistically significant correlations did exist at each DMR frequency, but that the frequency dependence was inconsistent with vibrational dust emission alone and strongly suggestive of additional dust-correlated free-free emission ($\\beta_{ff} = -2.15$).  A number of recent microwave observations have also shown that these correlations are not strongly dependent on angular scale. Different experiments ( DMR: Kogut et al. 1996b; 19GHz: de Oliviera-Costa et al. 1998; Saskatoon: de Oliviera-Costa et al. 1997; MAX5: Lim et al. 1996; OVRO: Leich et al. 1997), together give, with $95\\%$ confidence, $-3.6 < \\beta_{radio} < -1.3$, consistent with free-free emission over the frequency range 15-50~GHz.  However if the source of this correlated emission is indeed free-free, it is expected that a similar correlation should exist between $H_\\alpha$ and dust maps.  Several authors have found that these data sets are only marginally correlated (McCallough 1997 and Kogut 1997).  Thus, most of the correlated emission appears to come from another source.  Draine and Lazarian (1998a,b) suggest that the correlated emission could originate from spinning dust grains.  This model predicts a microwave emission spectrum that peaks at low microwave frequencies, the exact location of the peak depending on the size distribution of dust grains, and has a spectral index between -3.3 and -4, over the frequency range 15-50~GHz (Draine and Lazarian 1998b, Kogut 1999). The spectral indices obtained from various experiments do not seem to conform too well to the spinning dust model alone. However, recently de Oliveira-Costa {\\em et al} (1999) (hereafter DOC99) claim to have found evidence of the spinning dust origin of DIRBE-correlated emission using the Tenerife 10 and 15GHz data.  They find a rising spectrum from 10GHz to 15GHz, which is indicative of a spinning dust origin, as opposed to a free-free origin, for the dust-correlated Galactic foreground at these frequencies.  In this paper, we use the most recent Tenerife data to show that the correlated emission is not necessarily indicative of spinning dust. ",
        "conclusions": "Spinning dust emission can be identified and discriminated from free-free emission {\\em if} indications for a peak in the emission spectrum can be found at low microwave frequencies, which are probed by the Tenerife experiment. A rising spectrum between 10 and 15~GHz can be taken as a prediction of the spinning dust hypothesis, although the exact location of the emission peak depends on details of the model. Tentative evidence for a turnover in the spectrum of the Galactic dust-correlated microwave component between 10 and 15~GHz has been presented by DOC99. We however find that the spectral index of dust-correlated emission is negative for all Galactic cuts except for the $b > 20\\deg$ cut. The Galactic signal, and with it the significance of the correlation, decreases with increasing Galactic latitude, and no correlations are detected in the higher Galactic latitude regions of $b > 30\\deg$. The variance in the value of $\\hat{a}$ between regions with different Galactic cuts is rather large. We further find that the correlation detected in the $b> 20\\deg$ region comes only from a small number of pixels at low Galactic latitude and towards the Galactic centre, where signal from the Galactic plane is present. This correlation shows a free-free like spectral index, whereas the rest of the region was found to be uncorrelated, or even significantly anticorrelated.  Using sky rotations we show that the correlation we see at $b > 20\\deg$ is only an alignment of structure due to the rise of the Galactic plane signal. Employing a simple model for this structure we were able to demonstrate that the spatial distribution of Galactic emission is in fact different in the templates and in the data, giving rise to a systematic error. We were also able to show that this simple model of the Galaxy fits the data generally better than the templates. Further, modelling a Galactic free-free component, which correlates with the dust template, generally yields an equally acceptable fit to the data. Another significant systematic effect arises due to intrinsic errors in the templates, and all these effects cause a misleading increase in the inferred spectral index of the dust-correlated component between 10 and 15~GHz.  A comparison of our results with other experiments is presented in Figure 7.  Here the data points for the Tenerife experiment, obtained by taking a weighted average of all detections (all regions of all Galactic cuts, joint analyses) with errors taken from the sky rotations, correspond to values of $180\\pm47$ and $123\\pm16$ $\\mu$K /MJy sr$^{-1}$ at 10 and 15~GHz respectively. Note that the 10GHz point on the plot is significantly higher, as compared to that plotted in DOC99 and the value quoted in our table 2 This is because this region consists of parts that are correlated as well as parts that do not correlate or even anticorrelate, as in the case of the 10GHz data. Here, since we are focussing only on the regions that correlate (these regions are the same for both 10 and 15~GHz data and have been found to lie close to the Galactic plane), the value is significantly higher. Note that the high values of $a$ that we get need not be in contradiction to the typical level of free-free allowed by $H_{\\alpha}$ maps of other regions as our detections are close to the Galactic plane, where the level of Galactic emission is expected to be high.  \\begin{figure} \\centerline{\\epsfig{ file=spectra.eps,height=8cm,width=8cm,angle=0}} \\caption{Summary of our results on the dust-correlated (top panel) and synchrotron-correlated (bottom panel) emission compared to results from other experiments (COBE DMR-filled circles, 19GHz-empty triangle, OVRO-filled triangle, Saskatoon-filled square). The straight lines represent synchrotron (solid), free-free (dotted) and vibrational dust (dashed) spectral indices. The thick solid curve shows a combination of free-free emission with a small spinning dust contribution with a peak at 15 GHz in an attempt to fit all the data. This fit appears acceptable, as does the single free-free component fit, but it should be kept in mind that the data from different experiments might not be directly comparable, since they were taken on different angular scales and towards different regions of the sky. The Tenerife data points are an average derived from our various results, see text for details.} \\end{figure}  The spectral index for dust-correlated emission as deduced from the 10 and 15 GHz points is less negative (by about $1\\sigma$) than expected for free-free emission. The spectral index deduced simultaneously from synchrotron-correlated emission is systematically steeper than expected. This could be attributed to the effects that we have identified and which have all been shown to influence the correlation in the same direction.  Note also that in this plot only the Tenerife and COBE DMR points, at different frequencies, each represent data which probe the same angular scales and were taken with the same sky coverage. Inferring a spectral index by comparing different experiments assumes that the Galactic component is traced by a given fixed template and does not depend on parameters which vary over the sky, but our analysis of the Tenerife data shows large variations of Galactic correlation with the sky region.  With the present analysis of Tenerife data we are not able to make a firm claim about the origin of the dust-correlated component, since we do not find convincing support for a spinning dust component. This does not have to rule out this hypothesis, since environmental conditions or grain sizes, which affect the position of the spectral peak, could change systematically with location on the sky, particularly in the transition region between low and high Galactic latitude.  The spatial variation in the correlation amplitude and the possibility of the presence of some spinning dust emission along with free-free emission need to be dealt with. The separation task is difficult for two reasons.  Firstly neither the frequency dependence nor the spatial distribution of any of the Galactic components at low microwave frequencies is particularly well known. And further we would expect these components to be correlated, since we find their templates to be correlated, at least at low Galactic latitudes. The combination of the results from the other experiments, shown in Figure 7, might not be strongly constraining, but nevertheless, does not give conclusive evidence for spinning dust emission either, without an expectation of the amplitude of free-free emission based on dust-correlated H$_\\alpha$ emission.  In order to make more reliable inferences, a pixel by pixel separation of components would have to be performed, using the Maximum Entropy method for example.  In a forthcoming paper we shall perform such a separation incorporating information about spinning dust. Also, including other data at frequencies lower than 10~GHz and adding other templates of Galactic emission such as the $H_\\alpha$ maps from the Wisconsin H-Alpha mapper (Tufte, Reynolds and Haffner 1998) would be useful."
    },
    "0002/astro-ph0002133_arXiv.txt": {
        "abstract": "We present simple models for disk evolution based on two different approaches: a forward approach based on predictions generic to hierarchical models for structure formation (e.g., Mo, Mao, \\& White 1998) and a backwards approach based on detailed modeling of the Milky Way galaxy (e.g., Bouwens, Cay\\'on, \\& Silk 1997).  We normalize these models to local observations and predict high-redshift luminosities, sizes, circular velocities, and surface brightnesses. Both approaches yield somewhat similar predictions for size, surface brightness, and luminosity evolution though they clearly differ in the amount of number evolution.  These predictions seem to be broadly consistent with the high-redshift observations of Simard et al.\\ (1999), suggesting that the $B$-band surface brightness of disks has indeed evolved by $\\sim1.5^m$ from $z\\sim0$ to $z\\sim1$ similar to the models and is not an artifact of selection effects as previously claimed.  We also find a lack of low surface brightness galaxies in several high redshift samples relative to model predictions based on local samples (de Jong \\& van der Kruit 1994; Mathewson, Ford, \\& Buchhorn 1992). ",
        "introduction": "Over the last several years, there has been a steady increase in the number and quality of observations available for disk galaxies from $z=0$ and $z=1$.  Schade et al.\\ (1995,1996), using early ground and space based images of galaxies from the Canada-France Redshift Survey (CFRS), found a net increase in the surface brightness of galaxies to $z\\sim1$.  Along the same lines, Roche et al.\\ (1998), compiling 347 galaxies from the Medium Deep Survey and other surveys, concluded that disk galaxies had undergone a net evolution in surface brightness and a net devolution in size.  Lilly et al.\\ (1998), using structural parameters extracted from HST images of the combined CFRS and LDSS2 sample, concluded that there has been essentially no evolution in large disks out to $z\\sim1$.  As a preliminary effort as part of the DEEP survey, Vogt et al.\\ (1996,1997) found little evolution in the Tully-Fisher relationship ($<0.3^m$) out to $z\\sim1$.  More recently, these observations have been augmented by the DEEP sample with 197 galaxies from the Groth strip to $I<23.5$, 1.5 magnitudes deeper than the LDSS2-CFRS sample.  In a first paper, Simard et al.\\ (1998) concluded that there had been little evolution in the disk surface brightness distribution to $z\\sim1$ contrary to previous claims.  A number of different approaches have been proposed for making specific predictions about disk evolution.  Mo, Mao, \\& White (1998a) showed how the standard paradigm for hierarchical growth of structure combined with simple assumptions about angular momentum conservation led to simple scaling relationships for the change in disk properties as a function of redshift.  Other authors (Ferrini et al.\\ 1994; Prantzos \\& Aubert 1995; Prantzos \\& Silk 1998; Boissier \\& Prantzos 1999; Chiappini, Matteucci \\& Gratton 1997), taking more of a backwards approach to the problem, used detailed studies of the profiles of the Milky Way and other nearby galaxies to propose radially dependent models of star formation in disk galaxies, models which could be used to make detailed predictions about high-redshift disk evolution.  Already there have been a number of elegant studies in which both the backwards approach (Cay\\'on, Silk, \\& Charlot 1996; Bouwens, Cay\\'on, \\& Silk 1997; Roche et al.\\ 1998) and the forwards approach (Mao, Mo, \\& White 1998; Steinmetz \\& Navarro 1999; Contardo, Steinmetz, \\& Fritze-von Alvensleben 1998; van den Bosch 1998; Mo, Mao, \\& White 1998b) have been used to interpret the observations available for disk galaxies, mostly to $z\\sim1$.  Unfortunately, none of these studies considered the important effect that a large spread in surface brightness could have on the interpretation of these observations, particularly the potentially large fraction of low surface brightness galaxies.  In some studies, the surface brightness selection effects at low and high redshift were simply ignored, and in others, e.g., Roche et al.\\ (1998), the spread was limited to $0.3\\,\\textrm{mag/arcsec}^2$ about Freeman's law (Freeman 1970).  Clearly, given the apparent large numbers of low surface brightness galaxies seen locally, it is quite logical to wonder if these galaxies are detectable in current high redshift surveys.  Indeed, one might wonder whether these galaxies or the observed correlation between luminosity and surface brightness may have already affected the interpretation of high redshift observations.  In light of the recent claim by Simard et al.\\ (1999) that the apparent surface brightness evolution thus far inferred to $z\\sim1$ is completely due to surface brightness selection effects, such a study would seem to be especially timely.  Secondly, none of these studies directly compared the predictions of the forward and backward approaches using the same observations.  Simple comparisons of the scaling expected in surface brightness, size, luminosity, and number are useful for interpreting the high redshift observations.  To address these shortcomings, we shall therefore consider implementations of both approaches, normalize them to the observed $z\\sim0$ size-luminosity relationship, compare their predictions, and consider how each of them fares at explaining the observed disk evolution out to $z\\sim1$ incorporating all the selection effects as they are best understood.  We commence by presenting our models (\\S2) and the observational samples with which we compare (\\S3).  We present the results (\\S4), discuss them (\\S5), and then summarize our conclusions (\\S6).  Throughout this study, we use $H_0 = 50\\,\\textrm{km/s/Mpc}$ unless otherwise noted. ",
        "conclusions": "There is a real question about a lack of low surface brightness galaxies relative to our predictions, especially as compared to the no-evolution model predictions.  This conclusion is somewhat dependent on the assumed correlation between surface brightness and luminosity as is evident in Figure 4.  This conclusion is also dependent on the selection biases against low surface brightness galaxies not being stronger than those considered here.  There is an extensive literature discussing surface brightness selection biases (Disney 1976; Allen \\& Shu 1979) and various attempts to derive the bivariate luminosity-surface brightness distribution of galaxies (McGaugh 1996; Dalcanton et al.\\ 1997b; Sprayberry et al.\\ 1997).  Surface brightness has a particularly strong effect on isophotal magnitude determinations, especially for low surface brightness galaxies; and this can introduce significant errors in the magnitude determinations, so the effective volume probed for these galaxies is significantly smaller than it is for equivalent luminosity high surface brightness galaxies (McGaugh 1996).  Simard et al.\\ (1999) in a detailed quantification of the selection effects of the DEEP sample do not consider the effect of surface brightness on the magnitudes and sizes recovered since typical errors were found to be $0.2^m$ (Simard 1999, private communication). Despite the relatively small size of this error, it is not entirely clear to the present authors that the errors would not become quite significant for the lowest surface brightness galaxies in the sample, particularly those just marginally detectable given the chosen object identification and photometric parameters.  Secondly, Simard et al.\\ (1999) considers disk galaxies to be optically thin whereas the observations of Lilly et al.\\ (1998) are more consistent with disks being optically thick.  Highly-inclined optically thin disks would be much more detectable than face-on or optically-thick disks.  The upshot is that at many apparent magnitudes and radii, Simard et al.\\ (1998) would suppose that at least some highly inclined galaxies would be detectable and therefore the selection function $S_{UP}$ there would be non-zero when in reality if disks were optically thick it would be zero.  For these reasons, we used a slightly more conservative selection function in surface brightness than that given in Figure 4 of Simard et al.\\ (1999) (see \\S3.2).  Another possibility, not considered here, is that low surface brightness galaxies might form relatively late, meaning that their mass-to-light ratios remain relatively large until relatively recent epochs.  Of course, prima facie, this would seem unlikely given the apparently constant slope in the Tully-Fisher relationship to faint magnitudes.  In their own analysis of their sample of $\\sim200$ galaxies, Simard et al.\\ (1999) concluded that there had been little evolution in the surface brightness distribution of disk galaxies when all selection effects had been carefully considered, quite in contrast to our estimated $\\sim1.5^m$ of $B$-band surface brightness evolution. Little consideration, however, was paid to the evolution in the total \\textit{numbers} of high surface brightness galaxies.  Here, we find that the number of high surface brightness galaxies dramatically exceeds that predicted by the evolutionary models considered here, and we have argued that this provides evidence for an evolution in the surface brightness distribution of disk galaxies.  Our interpretation seems to be furthermore supported by the lack of low surface brightness galaxies relative to our models.  For no-evolution in the disk surface brightness distribution really to be present as Simard et al.\\ (1999) claims, high-redshift intervals should have similar numbers of low surface brightness galaxies to those found in local samples, and these galaxies seem to be deficient, even with respect to our models which show significant evolution in surface brightness.  While the conclusions of Simard et al.\\ (1999) appear to have been carefully drawn, we would like to suggest that there are significant uncertainties in their determination of the mean surface brightnesses in the lowest redshift intervals and therefore the inferred evolution in surface brightness due to the small size of the low redshift samples considered.  By applying the selection effects from the high-redshift bin identically to all redshift intervals, Simard et al.\\ (1999) restricted their analysis to that fraction of disk galaxies exceeding the high-redshift surface brightness detection limit.  Applying these selection criteria uniformly to all low redshift intervals severely pares down the low-redshift samples and significantly increases the uncertainty of their average surface brightness measure.  Given the observed range in observed surface brightness ($\\sim 2\\, \\textrm{mag/arcsec}^2$) and typical numbers ($\\sim 5-6$) for the lowest redshift bins, there is a non-negligible uncertainty in the average surface brightness at low redshift, $\\sim 0.6^m$.  Our estimates of $\\sim1.5^m$ of $B$-band surface brightness evolution are somewhat larger than that inferred by most authors.  Roche et al.\\ (1998) found $0.9^m$ of surface brightness evolution from $z\\sim0.2$ to $z\\sim0.9$, Lilly et al.\\ (1998) found $0.8^m$ of surface brightness evolution in their large disk sample, and Schade et al.\\ (1995,1996a) inferred $1.2^m$ and $1.5^m$ respectively to $z\\sim0.8$. Despite different differential measures of surface brightness evolution, most of these samples give similar values for the mean disk surface brightness near $z\\sim1$: $20.79 \\pm 0.17$ for the Roche et al.\\ (1998) sample ($0.65<z$), $19.9 \\pm 0.2$ for the Simard et al.\\ (1999) sample ($0.9<z<1.1$), $20.7 \\pm 0.25$ for the Lilly et al.\\ (1998) large disk sample ($0.5<z<0.75$), $20.2 \\pm 0.25$ for the Schade et al.\\ (1995) sample ($0.5<z$), and $19.8 \\pm 0.1$ for the Schade et al.\\ (1996) sample ($0.5<z<1.1$).  Consequently, differences in the surface brightness evolution inferred derive from differences in the $z=0$ surface brightness distributions assumed.  We assume a distribution consistent with the local data of de Jong \\& van der Kruit (1994) and Mathewson et al.\\ (1992) as a baseline for measuring evolution with a surface brightness peaking faintward of Freeman's Law ($\\sim 21.7 B\\,\\textrm{mag/arcsec}^2$; Freeman 1970) while Schade et al.\\ (1995,1996) simply makes reference to Freeman's Law ($\\sim 21.65\\,b_j\\,\\textrm{mag/arcsec}^2$).  Simard et al.\\ (1999), Lilly et al.\\ (1998), and Roche et al.\\ (1998) measure surface brightness evolution differentially across their samples.  Possible problems here are surface brightness selection effects and limited low-redshift samples."
    },
    "0002/astro-ph0002180_arXiv.txt": {
        "abstract": "We present spectrophotometry in the 3600--9700 \\AA\\ region for a sample of 39 \\hii\\ regions in the Galaxy and Magellanic Clouds, for which independent information is available on the spectral types and effective temperatures of the ionizing stars.  The spectra have been used to evaluate nebular diagnostics of stellar temperature, metal abundance, and ionization parameter, and compare the observed behavior of the line indices with predictions of nebular photoionization models.  We observe a strong degeneracy between forbidden-line sequences produced by changes in stellar \\teff\\/ and metal abundance, which severely complicates the application of many forbidden-line diagnostics to extragalactic \\hii\\ regions.  Our data confirm however that the Edmunds \\& Pagel [\\ion{O}{2}]$+$[\\ion{O}{3}] abundance index and the V\\'{\\i}lchez \\& Pagel \\etap\\ index provide more robust diagnostics of metal abundance and stellar effective temperature, respectively. A comparison of the fractional helium ionization of the \\hii\\ regions with stellar temperature confirms the reliability of the spectral type vs \\teff\\/ calibration for the relevant temperature range \\teff\\ $\\le$ 38000 K.  We use empirical relations between the nebular hardness indices and \\teff\\/ to reinvestigate the case for systematic variations in the stellar effective temperatures and the upper IMFs of massive stars in extragalactic \\hii\\ regions.  The data are consistent with a significant softening of the ionizing spectra (consistent with cooler stellar temperatures) with increasing metal abundance, especially for $Z \\le$ \\zsun.  However unresolved degeneracies between $Z$ and \\teff\\ still complicate the interpretation of this result. ",
        "introduction": "Over the past 15 years high-quality emission-line spectra have been obtained for hundreds of \\hii\\ regions in nearby spiral and irregular galaxies.  The primary use of these data has been to study the chemical composition patterns in galactic disks (e.g., McCall, Rybski, \\& Shields 1985, Vila-Costas \\& Edmunds 1992, Zaritsky, Kennicutt, \\& Huchra 1994, Dinerstein 1996).  However with the availability of state-of-the-art photoionization codes such as {\\sc cloudy} (Ferland \\etal 1998), combined with stellar population synthesis codes for young clusters (e.g., Leitherer \\& Heckman 1995, Leitherer et al. 1999), \\hii\\ region spectra are being applied increasingly to measure the ages and initial mass functions (IMFs) of the ionizing star clusters (e.g., V\\'{\\i}lchez \\& Pagel 1988, Shields 1990, Stasi\\'nska \\& Leitherer 1996, Bresolin, Kennicutt, \\& Garnett 1999, hereafter BKG). Similar techniques have been applied to model the infrared emission-line spectra of IR-luminous starbursts, and the results have been interpreted as evidence for an anomalous IMF in these objects (e.g., Rieke et al. 1993).  A basic limitation in this approach is its reliance on a long chain of theoretical inputs: stellar evolution models as functions of stellar mass, chemical composition and mass-loss rates; stellar atmosphere models as functions of effective temperature, surface gravity, chemical composition, and mass-loss properties, and photoionization models for the surrounding \\hii\\ region.  The derived nebular abundances tend to be insensitive to the details of these models, but the inferred properties of the ionizing stars are sensitive to the model inputs at every step in the chain. The problem is especially acute when using \\hii\\ region or starburst spectra to constrain the effective temperatures and IMFs of the ionizing stars (e.g., Mathis 1985, McCall et al.~1985, V\\'{\\i}lchez \\& Pagel 1988, BKG). Most studies show evidence for a systematic softening of the ionizing continuum of \\hii\\ regions with increasing metal abundance, but the absolute range in stellar effective temperatures is highly model dependent, and the case for a systematic variation in the upper IMF is shaky at best (BKG and references therein).  Another approach, which can circumvent many of these difficulties, is to test the nebular diagnostics empirically, by obtaining spectra for nearby \\hii\\ regions which are ionized by stars of known spectral type. This makes it possible to anchor the nebular indicators of stellar temperature directly to the stellar spectral classification system. Furthermore, by modeling these \\hii\\ regions with the same methods that are applied to extragalactic regions one can directly assess the reliability of the model-based approach.  This empirical approach was first explored by Chopinet \\& Lortet-Zuckermann (1976) and Kaler (1978) to calibrate the [\\ion{O}{3}]$\\lambda$5007/\\hbeta\\ ratio as an indicator of stellar effective temperature.  For studies of extragalactic \\hii\\ regions and starbursts a more robust hardness index than [\\ion{O}{3}]/\\hbeta\\ is desirable, because the excitation of [\\ion{O}{3}] varies locally within \\hii\\ regions, and it is sensitive to other physical parameters (e.g., metal abundance, ionization parameter, dust) which can vary systematically in galaxies (e.g., Kennicutt 1984, Shields 1990). Other hardness indices have been developed for extragalactic applications (e.g., V\\'{\\i}lchez \\& Pagel 1988), but their reliability has not been tested empirically.  In this paper we report the results of a spectrophotometric study of 39 \\hii\\ regions in the Galaxy, the LMC, and the SMC, which is aimed at providing an empirical foundation for nebular-based measurements of the stellar ionizing continuum and IMF in galaxies.  We use the spectra to evaluate several nebular diagnostics of massive stellar populations, and to compare the known properties of the stars in these regions with those inferred from photoionization modeling. Our study is mainly motivated by applications to stellar temperatures and IMFs, but the approach is more generally applicable to spectral modeling of extragalactic \\hii\\ regions and starbursts. This paper is complementary in approach to a recent analysis of Oey et al. (2000), which has addressed many of these same questions using detailed point-by-point observations of 4 \\hii\\ regions in the LMC.  The remainder of this paper is organized as follows. The \\hii\\/ region sample and the observations are described in \\S~2 and \\S~3, respectively.  In \\S~4 we analyze the behavior of the principal diagnostic line ratios and investigate several commonly applied abundance and spectral hardness indicators. In \\S~5 we apply the empirical calibration to constrain the range of stellar effective temperature and upper IMF in galactic disks. ",
        "conclusions": "In conclusion, we have presented a set of empirical diagnostic line ratios, calibrated against the effective temperature of the exciting stars. Our results provide a stronger empirical foundation for many of the most widely used nebular diagnostics (e.g., \\x, \\etap), and they show much better agreement between observations and models than often was seen in earlier generations of stellar atmosphere and nebular photoionization models (cf., Mathis 1985, V\\'ilchez \\& Pagel 1988). The empirical relations presented here should be useful in many investigations of the stellar populations of extragalactic \\hii\\/ regions, where very limited information is available on the exciting stars, and where an estimate of the nebular-based effective temperature is needed, even with a restricted wavelength coverage.  We close by reemphasizing the exploratory nature of this study, and pointing out the most urgently needed improvements before one can reliably extract quantitative constraints on the stellar contents and IMFs of star forming regions from nebular spectra. On the empirical side, this approach can be much improved if accurate ($T_e$-based) metal abundances are obtained for all of the calibrating \\hii\\ regions (or their exciting stars), and the effects of effective temperature, metal abundance, and nebular geometry are accounted for explicitly in the analysis.  Oey et al. (2000) have demonstrated the value of obtaining high signal/noise point-by-point measurements of individual \\hii\\ regions, and applying this approach to a larger sample of objects would provide a major improvement over our exploratory analysis.  On the theoretical side, more self-consistent stellar models for the relevant range of abundances, temperatures, and gravities are the most crucial inputs.  This would also allow for a consistent treatment of stellar and nebular abundances in the photoionization calculations.  We hope that these initial results may stimulate further progress in these areas."
    },
    "0002/astro-ph0002463_arXiv.txt": {
        "abstract": "We review the history and a few features of IC/3K emission.  We discuss observing strategies, the physical parameters which can be derived directly or indirectly from successful detections, and the disparities between magnetic field strength estimates found by different methods. ",
        "introduction": "The recent history of inverse Compton emission from cosmic background photons and relativistic electrons (IC/3K) began during the period when the origin of extended X-ray emission from clusters of galaxies was unclear.  Felten and Morrison (1966), Blumenthal \\& Gould (1970), and Jones, O'Dell \\& Stein (1974) were among those who worked out the details of IC/3K, but not until iron lines were detected in the X-ray spectra was it generally accepted that the bulk of the X-ray emission was caused by the thermal bremsstrahlung process.  Although there were many attempts to isolate IC/3K emission from clusters, these were thwarted by the strong thermal emission which is quasi ubiquitous in clusters.  The recent breakthrough came from the dual detections of IC/3K from the lobes of Fornax A: spatially with the ROSAT PSPC (Feigelson et al. 1995) and spectrally with ASCA (Kaneda et. al. 1995).  In the last two years, there have been papers published on other sources which are relatively free of hot gas emission and from cluster radio halos, both from EUV excesses and from the harder spectral component at energies for which the thermal X-rays are dropping exponentially. ",
        "conclusions": "We have a fairly good idea of the problems involved in using IC/3K emission to derive physical quantities.  Most of these difficulties are not going to go away soon, but the quality of the X-ray data should improve significantly with the advent of several missions with larger collecting area, better resolution, and extended frequency coverage."
    },
    "0002/astro-ph0002149_arXiv.txt": {
        "abstract": " ",
        "introduction": "A statement of statistical belief not uncommon in cosmic ray work is: \"you need five sigmas to convince me\". This has some justification, in that the history of cosmic rays contains many instances when a source or effect is claimed but not subsequently substantiated. Frequently this has been due to incorrect application of some statistical technique, often a failure to account fully  for the 'degrees of freedom'.  Most of the present body of statistical knowledge has been developed for specific problems, few of which occur in cosmic rays, although one of the most useful of texts\\cite{kn:eadie} was produced for experimental particle physicists. The analyser of cosmic ray data has particular problems: cosmic ray data requires great effort in collection and they are unlikely, once analysed, to be repeated. The numbers are frequently small and there are usually data missing and frequently there is significant contamination by noise. Ideally, a statistical measure should be developed specifically for each application. This is the only way in which all of the parameters of the experiment can be allowed for in the analysis. It is more usual for a general statistical tool to be applied, for example \\chisq, which may not be optimal for the purpose, and for some experimental variables to be ignored.  The focus of this review will be on methods of determining the presence of a signal rather than estimating some parameters of the data. The aim is to gather together the recent developments in methods of analysis of the temporal and spatial features of cosmic ray data, especially where the methods used are not 'traditional'.  Several new methods have been published recently which depend on Bayesian ideas, and these ideas have been introduced before the description of the methods. ",
        "conclusions": "The hope of this review is that the more commonly met data analysis problems may be approached by the cosmic ray worker with a more consistent and up to date approach. There have been a number of advances in recent years in the tools, and more importantly in the methods, available to cosmic ray experimenters to ensure that the maximum use is made of hard-won data. The traditional statistical methods have resulted in a measure of agreement on the 'correct' way to look for sources from ON/OFF data, change points (bursts) in 1- and 2-dimensions and in periodicity. The application of these methods requires care to ensure that the 'degrees of freedom' are kept under control and properly accounted for: many of the criticisms of claimed sources have been based on the latter.  New Bayesian methods of testing hypotheses have recently been proposed. A central theme of these methods is that classical methods often cloak ignorance in a way which distorts the results. There are claimed to be significant benefits to the use of Bayesian methods which derive from the requirement to be absolutely specific about the hypotheses and the methodology of marginalising nuisance parameters. In contrast to classical statistical methods, where various \\emph{statistics} may be generated from the same data, each with different assumptions, \\emph{degrees of freedom} and {power}, Bayesian methods provide a framework for describing completely the data and allow the direct comparison of specified hypotheses. A practical result of the philosophical differences between the approaches is that, rather than relying on a relatively easy-to-use, pre-packaged test statistic, with the accompanying dangers of hidden degrees of freedom, a Bayesian method requires the data interpreter to model the hypotheses precisely. The obvious disadvantages of this are claimed to be more than compensated by the directness of the link between the hypotheses and the data. Bayesian methods may require some time to become accepted in the field, in that the methodologies and ideas have not traditionally been part of the training of physicists; indeed may not have been as useful if physicists' training in classical methods had been better."
    },
    "0002/astro-ph0002239_arXiv.txt": {
        "abstract": "We report the serendipitous discovery of two objects close in projection with fairly strong emission lines at long wavelength ($\\lambda\\sim 9190$\\,\\AA). One (A) seems not to be hosted by any galaxy brighter than $V_{555}=27.5$, or $I_{814}=26.7$ (Vega-based 3$\\sigma$ limits in 1\\farcs0 diameter apertures), while the other line is associated with a faint ($I_{814}\\simeq 24.4$) red galaxy (B) offset by 2\\farcs7 and 7 \\AA\\ spectrally.  Both lines are broad (FWHM $\\approx 700$\\,km\\,s$^{-1}$), extended spatially, and have high equivalent widths ($W_\\lambda^{\\rm obs}({\\rm A}) > 1225$ \\AA, 95\\% confidence limit; $W_\\lambda^{\\rm obs}({\\rm B}) \\approx 150$ \\AA).  No secondary spectral features are detected for galaxy A.  Blue continuum and the marginal detection of a second weak line in the spectrum of galaxy B is consistent with \\oxytwo\\ (the strong line) and \\magtwo\\ (the weak line) at $z = 1.466$.  By association, galaxy A is likely at $z=1.464$, implying a rest-frame equivalent width of the \\oxytwo\\ emission line in excess of 600\\,\\AA\\ and a projected separation of 30~$h_{50}^{-1}$~kpc for the galaxy pair.  Conventional wisdom states that isolated emission lines with rest-frame equivalent widths larger than $\\sim 200$ \\AA\\, are almost exclusively \\lya.  This moderate-redshift discovery therefore compromises recent claims of high-redshift \\lya\\ emitters for which other criteria (\\ie line profile, associated continuum decrements) are not reported.  We discuss observational tests to distinguish \\lya\\ emitters at high redshift from foreground systems. ",
        "introduction": "Serendipitous detections of emission line galaxies are common on low-dispersion spectrograms taken with large ground-based telescopes. Indeed, finding distant galaxies through blank sky slit spectroscopy is fully complementary to narrow-band imaging searches for distant line-emitting galaxies:  rather than probing a large area of sky for objects over a limited range of redshift, deep slit spectroscopy surveys a smaller area of sky for objects at a larger range in redshift \\markcite{Pritchet:94, Thompson:95}(\\eg Pritchet 1994; Thompson \\& Djorgovski 1995).  Also, since the resolution of optical spectra is better matched to narrow line emission than filters with widths of $\\approx 3000$ \\kms, deep slit spectroscopy is vastly more sensitive than narrow-band imaging.  In a 1.5 hour spectrum with the Keck telescope at moderate dispersion ($\\lambda / \\Delta \\lambda \\simeq 1000$), the limiting flux density probed for spectrally unresolved line emission in a 1 arcsec$^2$ aperture is $S_{\\rm lim} (3 \\sigma) \\approx 6 \\times 10^{-18} \\ergcm2s$ at $\\lambda \\approx 9300$ \\AA.  The first galaxy confirmed at $z > 5$ was found serendipitously during spectroscopic observations of a galaxy at $z = 4.02$ \\markcite{Dey:98}(Dey {et~al.} 1998).  Since then, several other $z > 5$ galaxies have been reported \\markcite{Hu:98, Weymann:98, Spinrad:98, vanBreugel:99a, Chen:99, Hu:99}(Hu, Cowie, \\& McMahon 1998; Weymann {et~al.} 1998; Spinrad {et~al.} 1998; van Breugel {et~al.} 1999; Chen, Lanzetta, \\& Pascarelle 1999; Hu, McMahon, \\& Cowie 1999).  In both narrow-band imaging and slit spectroscopy surveys, the equivalent width of detected emission lines is a standard redshift indicator \\markcite{Cowie:98, Hu:98}(\\eg Cowie \\& Hu 1998; Hu {et~al.} 1998).  \\lya\\ at high redshift can have a very large rest-frame equivalent width \\markcite{Charlot:93}(up to $W_\\lambda^{\\rm rest} \\approx 200$ \\AA\\ if driven by star formation; Charlot \\& Fall 1993), while \\oxytwo $\\lambda$3727\\AA\\ at low to moderate redshift, the other primary strong, solitary emission feature in UV/optical galaxy spectra, rarely has a rest-frame equivalent width exceeding 100 \\AA\\, in magnitude-limited surveys \\markcite{Songaila:94, Guzman:97, Hammer:97, Hogg:98}(\\eg Songaila {et~al.} 1994; Guzm\\`an {et~al.} 1997; Hammer {et~al.} 1997; Hogg {et~al.} 1998).  Equivalent width selection is also helped by $W_\\lambda^{\\rm obs} = W_\\lambda^{\\rm rest} (1 + z)$, which boosts \\lya\\ more than \\oxytwo.  The other strong UV/optical emission features in galaxies, \\eg \\hbeta, \\oxythree$\\lambda\\lambda$4959,5007\\AA, \\halpha, are generally easily identified spectroscopically from their wavelength proximity to other emission features, though \\halpha\\ can also have extremely high equivalent widths, up to 3000 \\AA\\ \\markcite{Leitherer:95}(Leitherer, Carmelle, \\&  Heckman 1995), leading to some confusion between low-redshift, young, dwarf starbursts and high-redshift \\lya-emitters \\markcite{Stockton:98}(\\eg Stockton \\& Ridgway 1998).  Here we report the discovery of two emission lines in the approximate direction of the Abell~2390 cluster.  The equivalent widths of these lines are rather large, with $W_\\lambda^{\\rm obs} (\\rm A) > 1225$ \\AA\\ (95\\% confidence limit) for the first object and $W_\\lambda^{\\rm obs} (\\rm B) \\approx 150$ \\AA\\, for the second object.  However, as we argue below, \\oxytwo\\ at $z=1.46$ is the most likely identification for these features, indicating an atypical system.  In \\S 2 $-$ \\S 4 we discuss the observations, redshift determination, and properties of these galaxies.  In \\S 5 we consider the implications of this discovery for narrow-band searches for \\lya\\ emission from distant protogalaxies, including a detailed discussion of the observational criteria useful for distinguishing high-redshift \\lya\\ from foreground emission-line galaxies.  Throughout this paper, unless otherwise indicated, we adopt $H_0 = 50~ h_{50}~ \\kmsMpc, q_0 = 0.1$, and $\\Lambda = 0$.  For these parameters, the luminosity distance, $d_L$, at $z=1.46$ is 13.88 $h_{50}^{-1}$ Gpc and 1\\arcsec\\ subtends 11.1 $h_{50}^{-1}$ kpc.    \\begin{figure}[t!] \\plotfiddle{stern.fig1.eps}{2.0in}{0}{91}{91}{-220}{0}  \\caption{{\\it HST} $I$-band (F814W; 10500 s) image of the serendipitously discovered galaxies in the field of Abell~2390.  Galaxy A, marked by dotted circle, is undetected ($I_{814} > 26.7, 3\\sigma$). Galaxy B has $I_{814} = 24.4$ (1\\arcsec\\ diameter aperture).  Slit positions for our three observations are indicated.  The images shown are 30\\arcsec\\ on a side, oriented with north upwards and east to the left.  Galaxy B is located at $\\alpha = 21^h53^m35\\secper16$, $\\delta = +17\\deg42\\arcmin56\\farcs66$ (J2000) and is 2\\farcs7 SE of galaxy A. The bright elliptical galaxy south of galaxy B is a member of the Abell 2390 galaxy cluster ($z=0.241$).}  \\label{hstslit} \\end{figure} ",
        "conclusions": "We report the serendipitous discovery of two faint galaxies with high equivalent width emission lines at long wavelength ($\\lambda \\approx 9190$ \\AA).  For one source, galaxy B, faint blue continuum and a weak secondary line emission feature implies the source is at $z = 1.466$. The spatial proximity and similar emission wavelength of object A persuasively argues that this is an unusual \\oxytwo-emitter at $z = 1.464$.  It had been thought that serendipitous one-lined sources with rest-frame equivalent widths larger than a few $\\times$ 100 \\AA\\ are exclusively identified with high-redshift \\lya.  Our observations have shown that this is demonstrably not the case.  Both sources are unlike local star-forming galaxies, and we suggest that the observations are most consistent with the discovery of a moderate-redshift active system.  Several programs are currently underway to find high-redshift \\lya\\ emitters using a combination of narrow- and broad-band imaging \\markcite{Thommes:98, Hu:98}(\\eg Thommes {et~al.} 1998; Hu {et~al.} 1998) and serendipitous long slit surveys \\markcite{Manning:00}(\\eg Manning {et~al.} 2000).  The moderate-redshift system discussed herein has serious implications for those surveys.  In particular we expect the line-emission surveys to uncover a population of strong line emitters and we emphasize that comparison of this sample with magnitude-limited surveys could be misleading.  We discuss various criteria which can be used to assess whether a solitary, high-equivalent width emission feature is associated with high-redshift \\lya.  In Table~2 we consider how recently reported $z > 5$ spectroscopic candidates fare with respect to these criteria.  Some sources are reliably confirmed at $z > 5$, while others are confirmed at low-redshift and yet others remain ambiguous with the current data. We conclude that some criteria are necessary, but not sufficient, to conclude that a source is at $z > 5$, such as a continuum decrement across the emission feature, while other criteria are sufficient, but not necessary, such as an asymmetric line profile.  We suggest that multiple criteria are necessary to convincingly demonstrate that a single-lined source is at high-redshift."
    },
    "0002/astro-ph0002525_arXiv.txt": {
        "abstract": "The two--fluid effects on the radial outflow of relativistic electron--positron plasma are considered. It is shown that for large enough Michel (1969) magnetization parameter $\\sigma \\gg 1$ and multiplication parameter $\\lambda=n/n_{GJ} \\gg 1$ one--fluid MHD approximation remains correct in the whole region $|{\\bmath E}| < |{\\bmath B}|$. In the case when the longitudinal electric current is smaller than the Goldreich--Julian one, % the acceleration of particles near the light surface $|{\\bmath E}| = |{\\bmath B}|$ is determined. It is shown that, as in the previously considered (Beskin Gurevich \\& Istomin 1983) cylindrical geometry, almost all electromagnetic energy is transformed into the energy of particles in the narrow boundary layer $\\Delta\\varpi/\\varpi \\sim \\lambda^{-1}$. ",
        "introduction": "Despite the fact that the structure of the magnetosphere of radio pulsars remains one of the fundamental astrophysical problems, the common view on the key theoretical question -- what is the physical nature of the neutron star braking -- is absent (Michel 1991, Beskin Gurevich \\& Istomin 1993, Mestel 1999). Nevertheless, very extensive theoretical studies in the seventies and the eighties allowed to obtain some model-independent results. One of them is the absence of magnetodipole energy loss. This result was first obtained theoretically (Henriksen \\& Norton 1975, Beskin et al 1983). It was shown that the electric charges filling the magnetosphere screen fully the magnetodipole radiation of a neutron star for an arbitrary inclination angle $\\chi$ between the rotational and magnetic axes if there are no longitudinal currents flowing in the magnetosphere. Later this result was also confirmed by observations. The direct detections of the interaction of the pulsar wind with a  companion star in close binaries (see e.g. Djorgovsky \\& Evans 1988, Kulkarni \\& Hester 1988) have shown that it is impossible to explain the heating of the companion by a low--frequency magnetodipole wave.  On the other hand, the detailed mechanism of particle acceleration remains unclear. Indeed, a very high magnetization parameter $\\sigma$ (Michel 1969) in the pulsar magnetosphere demonstrates that within the light cylinder $r < R_{\\rm L} = c/\\Omega$ the main part of the energy is transported by the Poynting flux. It means that the additional mechanism of particle acceleration must work in the vicinity of the light cylinder. It is necessary to stress that % an effective particle acceleration can only take place for small enough longitudinal electric currents $I < I_{GJ}$ when the plasma has no possibility to pass smoothly through the fast magnetosonic surface and when the light surface $|{\\bmath E}| = |{\\bmath B}|$ is located at a finite distance. As to the case of the large longitudinal currents $I > I_{GJ}$, both analytical (Tomimatsu 1994, Begelman \\& Li 1994, Beskin et al 1998) and numerical (Bogovalov 1997) considerations demonstrate that the acceleration becomes ineffective outside the fast magnetosonic surface, and the particle-to-Poynting flux ratio remains small: $\\sim \\sigma^{-2/3}$ (Michel 1969, Okamoto 1978).  The acceleration of an electron--positron plasma near the light surface was considered by Beskin Gurevich and Istomin (1983) in the simple $1D$ cylindrical geometry for $I \\ll I_{GJ}$. It was shown that in a narrow boundary layer $\\Delta\\varpi/\\varpi \\sim 1/\\lambda$ almost all electromagnetic energy is actually converted to the particles energy. Nevertheless, cylindrical geometry does not provide the complete  picture of particle acceleration. In particular, it was impossible to include self--consistently the disturbance of a poloidal magnetic field and an electric potential, the later playing the main role in the problem of the plasma acceleration (for more details see e.g. Mestel \\& Shibata 1994). Hence, a more careful $2D$ consideration is necessary.  In Sect. 2 we formulate a complete system of $2D$ two--fluid MHD equations describing the electron--positron outflow from a magnetized body with a monopole magnetic field. The presence of an exact analytical force--free solution (Michel 1973) allows us to linearize this system which results in the existence of invariants (energy and angular momentum) along unperturbed monopole field lines similar to the ideal one--fluid MHD flow. In Sect. 3 it is  shown that for $\\sigma \\gg 1$ and $\\lambda \\gg 1$ ($\\lambda=n/n_{GJ}$ is the multiplication factor) the one--fluid MHD approximation remains true in the entire region within the light surface. Finally, in Sect. 4 the acceleration of particles near the light surface $|{\\bmath E}| = |{\\bmath B}|$ is considered. It is shown that, as in the case of cylindrical geometry, in a narrow boundary layer $\\Delta\\varpi/\\varpi \\sim \\lambda^{-1}$ almost all the electromagnetic energy is converted into the energy of particles. ",
        "conclusions": "Thus, on a simple example it was demonstrated that for real physical parameters of the magnetosphere of radio pulsars ($\\sigma \\gg 1$ and $\\lambda \\gg 1$) the one--fluid MHD approximation remains true in the whole region within the light surface $|{\\bmath E}| = |{\\bmath B}|$. On the other hand, it was shown that in a more realistic 2D case the main properties of the boundary layer near the light surface existing for small enough longitudinal currents $I < I_{GJ}$ (effective energy transformation from electromagnetic field to particles, current closure in this region, smallness of the disturbance of electric potential and poloidal magnetic field) remain the same as in the 1D case considered previously (Beskin et al 1983).  It is necessary to stress the main astrophysical consequences of our results. First of all, the presence of such a boundary layer explains the effective energy transformation of electromagnetic energy into the energy of particles. As was already stressed, now the existence of such an acceleration is confirmed by observations of close binaries containing radio pulsars (as to the particle acceleration far from a neutron star, see e.g. Kennel \\& Coroniti 1984, Hoshino et al 1992, Gallant \\& Arons 1994). Simultaneously, it allows us to understand the current closure in the pulsar magnetosphere. Finally, particle acceleration results in the additional mechanism of high--energy radiation from the boundary of the magnetosphere (for more details see Beskin et al 1993).  Nevertheless, it is clear that the results obtained do not solve the whole  pulsar wind problem. Indeed, as in the cylindrical case, it is impossible to describe the particle motion outside the light surface. The point is that, as one can see directly from Eqn. (\\ref{o2}), for a complete conversion of electromagnetic energy into the energy of particles it is enough for them to pass only $\\lambda^{-1}$ of the total potential drop between pulsar magnetosphere and infinity. It means that the electron--positron wind propagating to infinity has to pass the potential drop which is much larger than their energy. It is possible only in the presence of electromagnetic waves even in an axisymmetric magnetosphere which is stationary near the origin. Clearly, such a flow cannot be considered even within the two--fluid approximation.  In our opinion, it is only a numerical consideration that can solve the problem completely and determine, in particular, the energy spectrum of particles and the structure of the pulsar wind.  Unfortunately, up to now such numerical calculations are absent."
    },
    "0002/hep-ph0002226_arXiv.txt": {
        "abstract": "We show that if the neutralino in the minimal supersymmetric standard model is the dark matter in our galaxy, there cannot be a dark matter cusp extending to the galactic center. Conversely, if a dark matter cusp extends to the galactic center, the neutralino cannot be the dark matter in our galaxy.  We obtain these results considering the synchrotron emission from neutralino annihilations around the black hole at the galactic center. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/hep-ph0002010_arXiv.txt": {
        "abstract": "We investigate the weak interaction emission of spin-$1/2$ fermions  from accelerated currents. As particular applications, we analyze the decay of uniformly accelerated protons and neutrons, and the neutrino-antineutrino emission from uniformly accelerated electrons. The possible relevance of our results to astrophysics is also discussed. ",
        "introduction": "\\label{Intro}   We investigate the weak interaction emission of spin-$1/2$ fermions from classical and semiclassical currents.  We denote by semiclassical those currents which possess classical trajectories and are endowed with quantized inner energy levels. Our results can be used to investigate a broad class of processes involving accelerated particles provided that they have a well defined world line, as indeed verified in many situations of interest. In the case where the particle is accelerated by a background electromagnetic field, such processes could be fully analyzed quantum-mechanically. As a consequence, any recoil effects  due to the fermion emission would be  automatically taken into account~\\cite{BLP,BKS}. For instance, in Ref.~\\cite{BLP} a quasiclassical approach to Quantum Electrodynamics was developed to consider $\\gamma$-synchrotron radiation from an electron in a classical background magnetic field. This  approach is basically characterized by assuming that the electron motion is quasiclassical. This is always possible as far as the magnetic field is not very strong, namely $H \\ll H_0$ and $\\gamma \\gg 1$, where $H_0=4.4 \\, \\cdot \\, 10^{13} $~Gauss and $\\gamma$ is the  Lorentz factor for the electron. This quasiclassical approach applied for neutrino-antineutrino emission is analyzed in detail in Sec.~6.1 of Ref.~\\cite{BKS}.  Although this is hard to take into account the current recoil in our semiclassical approach, the relations here obtained, which agrees with the full quantum mechanical treatment in the proper limit ($\\chi \\ll 1$, see e.g. Sec.~\\ref{Neutrino emission}), are easily applicable  when the process involves particle decay and the trajectory itself (rather than the underlying dynamical process which generates it) is  inferred from the observational data. Explicit results for uniformly accelerated currents are exhibited.   As far as we know, the first ones to call attention to the possibility that non-inertial protons may decay were Ginzburg and Syrovatskii~\\cite{GINZ} but only recently Muller~\\cite{M} obtained the first estimation for the decay rate associated with the process $ p \\to n \\;e^+  \\nu_e$ by assuming that all the involved particles are scalars. Here, as a particular application of modeling accelerated particles by  semiclassical currents, we perform a comprehensive (inertial-frame) analysis of the inverse  $\\beta$ decay for uniformly accelerated protons. We show that under certain astrophysical conditions, high-energy protons in strong background magnetic fields should rapidly decay.  The observation of non-inertial neutrons is less trivial. Notwithstanding the calculation of the $\\beta$ decay rate for accelerated neutrons may be of some relevance in situations where they are under the influence of ``relatively'' strong background gravitational fields and, thus, will be also performed.  Some features of the $\\beta$ and inverse  $\\beta$ decays for uniformly accelerated nucleons will be  discussed in terms of the  Fulling-Davies-Unruh (FDU) effect. The FDU  effect asserts that the Minkowski vacuum corresponds to a thermal bath with respect  to uniformly accelerated observers \\cite{FD,U}. It is perhaps remarkable that although inertial observers associate the inverse $\\beta$ decay to the channel $ p \\to n \\;e^+  \\nu_e$, coaccelerated observers associate the {\\em same} proton decay event to one of the following channels: $ p \\; e^- \\to n \\; \\nu_e , $ $ p \\; {\\bar \\nu}_e \\to n \\; e^+ $ or $ p \\; e^- {\\bar \\nu}_e \\to n\\; , $ where the absorbed $ e^-$ and ${\\bar \\nu}_e$ are Rindler particles present in the FDU thermal bath ``attached'' to the  proton's frame~\\cite{MV}. The corresponding branching ratios can be also calculated as a function of the proton acceleration.  Under a certain restriction, we can make our semiclassical current to behave as a classical one. This is suitable to investigate neutrino-antineutrino emission from accelerated electrons. This process is of relevance in some astrophysical situations as, e.g., in the cooling of neutron stars. Eventually we compare our results in the proper limit with the ones in the literature obtained by quantizing electrons in a background magnetic field \\cite{LP}-\\cite{people}.  The paper is  organized as follows: In Sec.~\\ref{DNC} we introduce the semiclassical currents and discuss how they model particle decays. In Sec.~\\ref{inertial}, we introduce the weak-interaction action and couple our current to a spin-$1/2$ fermion-antifermion field. Afterwards we calculate  the differential transition probability for currents following arbitrary world lines. In Sec.~\\ref{accelerated} we use the results obtained in the previous section to explicitly evaluate the fermion emission rate and radiated power for the particular case of a uniformly accelerated current. The next two sections, Sec.~\\ref{decay} and Sec.~\\ref{Neutrino emission}, are dedicated to analyze in detail the decay of uniformly accelerated protons and neutrons, and the neutrino-antineutrino emission from uniformly accelerated electrons, respectively. We also comment on the possible astrophysical relevance of our results. We dedicate Sec.~\\ref{Discussions} for our final discussions. We will use natural units $c = \\hbar = k_B = 1$ throughout this paper unless stated otherwise. ",
        "conclusions": "\\label{Discussions}   We have investigated the weak interaction emission of spin-$1/2$ fermions from classical and semiclassical currents. As a particular application of modeling the accelerated particle by a semiclassical current, we have analyzed the inverse $\\beta$ decay of uniformly accelerated protons. We have shown that although protons in laboratory storage rings are not likely to decay in this way, under some astrophysical conditions high-energy protons in background magnetic fields may have a considerably short lifetime. Moreover, we have analyzed the modification of the usual $\\beta$ decay for uniformly accelerated neutrons. This may be of some relevance when neutrons are under the influence of strong background gravitational fields. Although a full curved spacetime calculation is desirable to treat these situations, our calculation should be a good approximation when the gravitational field is ``moderate''~\\cite{CHM2}. In this case, neutrons  can be treated  as being accelerated in Minkowski space.  By restricting our semiclassical current to behave classically, we were able to use our formalism to investigate the neutrino-antineutrino pair emission from uniformly accelerated electrons and compare our results with the ones in the literature obtained by quantizing the electron field in a background magnetic field. Our formalism allows the utilization of currents associated with more general world lines. Depending on the accuracy level required, however, one can use directly the formulas derived for uniformly accelerated currents. This may be particularly useful in some astrophysical situations.      \\begin{flushleft} {\\bf{\\large Acknowledgements}} \\end{flushleft}  The authors are thankful to J.C. Montero, V. Pleitez and A.A. Natale for discussions. D.V. was fully supported by Funda\\c c\\~ao de Amparo \\`a Pesquisa do Estado de S\\~ao Paulo while G.M. was  partially supported by Conselho Nacional de Desenvolvimento Cient\\'\\i fico e Tecnol\\'ogico.  \\appendix"
    },
    "0002/hep-ex0002049_arXiv.txt": {
        "abstract": "The proton spectrum in the kinetic energy range 0.1 to 200\\,\\GeV\\ was measured by the Alpha Magnetic Spectrometer (AMS) during space shuttle flight STS--91 at an altitude of 380\\,km.  Above the geomagnetic cutoff the observed spectrum is parameterized by a power law.  Below the geomagnetic cutoff a substantial second spectrum was observed concentrated at equatorial latitudes with a flux $\\sim$70\\,m$^{-2}$sec$^{-1}$sr$^{-1}$. Most of these second spectrum protons follow a complicated trajectory and originate from a restricted geographic region. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002329_arXiv.txt": {
        "abstract": "We present the results of an analysis of the large scale velocity fields of the ionized gas associated with powerful radio galaxies. Long-slit spectra of 52 objects provide a sample of resolved velocities that span a wide range of redshifts, radio and emission-line luminosities. Line widths reaching 1000 km s$^{-1}$ and resolved velocity fields with amplitudes of up 1500 km s$^{-1}$ are found on scales from 10 to 100 kpc in the environments of radio galaxies at redshifts larger than 0.5. The global velocities and FWHM are of comparable amplitudes in the FRII sources, while the FRI sources have FWHM values that are larger than their resolved velocity fields.  We find evidence for systematically larger line widths and velocity field amplitudes at $z > 0.6$.  Several of the largest amplitude systems contain two galaxies with small projected separations. All of the $> 1000$ km s$^{-1}$ systems occur in objects at $z > 0.6$ and all have comparable radio and [OII] sizes. There is a weak correlation of off-nuclear line widths and velocity field with the ratio of the radio and emission-line sizes, but it is of low statistical significance and there is a very large dispersion. The change in properties at redshifts above $z \\sim 0.6$ could reflect a difference in environments of the host galaxies, with the hosts inhabiting higher density regions with increasing redshift (e.g., Hill \\& Lilly 1991). The mass of ionized gas and the apparent enclosed dynamical mass are correllated and both increase steeply with redshift and/or radio power.   The origin of the velocities remains uncertain.  The data do not require jet-gas interactions to explain the kinematics and superficially are slightly more consistent with gravitational origins for the bulk of the kinematics.  If the line width reflects the underlying gravitational potential, the observed FWHM traces the velocity dispersion of the host galaxy or its surrounding group or cluster. The highest velocities seen then point to interesting environments for intermediate and high redshift radio galaxies.  Turbulent interactions with the expanding radio source as the origin of the kinematics are certainly not ruled out. In the jet interaction scenario, the maximum velocities seen in the nebula can be used to constrain the density of the pre-shock gas to be roughly n$_{\\rm e} > 0.6$ cm$^{-3}$. ",
        "introduction": "We report the results of a study of the kinematics and morphology of the emission line nebula in powerful 3CR radio galaxies from redshift $\\sim$ zero to $\\sim 4$ based on the data presented in Baum, Heckman and van Breugel (1990) and McCarthy, Baum, and Spinrad (1996) and additional data gathered on these and similar sources from the more general literature. The 3CR is a well-defined radio-flux limited sample; radio power (AGN luminosity) is strongly correlated with redshift in this sample.  We investigate the correlation of the kinematic properties of the gas with other inherent properties, such as redshift, radio and line luminosity, radio and emission line morphology and extent. We also investigate correlations with derived properties such as inferred mass of emission line gas and apparent dynamical mass, where the apparent dynamical mass is defined to be the mass required to gravitationally induce the observed kinematics. We find that both the inferred mass in emission line gas and the apparent dynamical mass appear to increase with redshift (or radio power).  For each source we define a maximum emission line velocity (V$_{\\rm max}$, the maximum difference between gas velocity and the systemic velocity of the galaxy, measured as the emission line velocity at the continuum peak) and the projected radius at which that maximum velocity occurs (R$_{\\rm Vmax}$).  We find that across the full sample, the maximum emission line velocity increases with R$_{\\rm Vmax}$ and with the absolute size of the nebula. We also find that the sources with the highest maximum velocities (1000-1500 km sec$^{-1}$) are those with two independent kinematic systems present.  No correlations are seen between the size of the radio source and the FWHM or maximum velocity of the gas. Similar to the above, no substantial correlation  is seen between the FWHM or the maximum velocity of the radio source and the ratio of the radio to nebular size. Lastly, our sample shows no correlation of radio size with radio power or redshift, or of radio size with nebular size, but does show a substantial correlation of nebular size with redshift and/or radio power.  These results are consistent with (but do not necessitate) a picture in which the kinematics of the emission line gas are predominantly dictated by gravity and the line luminosity of the nebula is predominantly determined by the ionizing luminosity of the central engine.  If the mass of cold gas increases with redshift, then the increasingly powerful AGN present in the 3CR sample at higher redshift, will illuminate a larger ionized nebula. If the velocities are predominantly gravitational, then the most powerful sources which are at the highest redshifts in our sample, and show the highest emission line gas velocity, would have to be found in very deep gravitational potentials, such as cluster environments. Alternatively, at the highest redshifts, powerful radio galaxies may be in environments which are in a relatively young evolutionary state where the kinematics are still dominated by infall. In the gravitational interpretation, the very highest velocity systems, with apparently disjoint multiple velocity systems in the nebula of a given source, are cases where two galaxies are merging or interacting, and gas systems associated with each galaxy are seen. ",
        "conclusions": "We report the results of a study of the kinematics and morphology of the emission line nebula in powerful 3CR radio galaxies from redshift $\\sim$ zero to $\\sim 4$ based on the data presented in Baum, Heckman and van Breugel (1990) and McCarthy, Baum, and Spinrad (1996) and additional data gathered on these and similar sources from the more general literature. The 3CR is a well-defined radio-flux limited sample; radio power (AGN luminosity) is strongly correlated with redshift in this sample.  We investigate the correlation of the kinematic properties of the gas with other inherent properties, such as redshift, radio and line luminosity, radio and emission line morphology and extent. We also investigate correlations with derived properties such as inferred mass of emission line gas and apparent dynamical mass, where the apparent dynamical mass is defined to be the mass required to gravitationally induce the observed kinematics. We find that both the inferred mass in emission line gas and the apparent dynamical mass appear to increase with redshift (or radio power).  For each source we define a maximum emission line velocity (V$_{\\rm max}$, the maximum difference between gas velocity and the systemic velocity of the galaxy, measured as the emission line velocity at the continuum peak) and the projected radius at which that maximum velocity occurs (R$_{\\rm Vmax}$).  We find that across the full sample, the maximum emission line velocity increases with R$_{\\rm Vmax}$ and with the absolute size of the nebula. We also find that the sources with the highest maximum velocities (1000-1500 km sec$^{-1}$) are those with two independent kinematic systems present.  No correlations are seen between the size of the radio source and the FWHM or maximum velocity of the gas. Similar to the above, no substantial correlation  is seen between the FWHM or the maximum velocity of the radio source and the ratio of the radio to nebular size. Lastly, our sample shows no correlation of radio size with radio power or redshift, or of radio size with nebular size, but does show a substantial correlation of nebular size with redshift and/or radio power.  These results are consistent with (but do not necessitate) a picture in which the kinematics of the emission line gas are predominantly dictated by gravity and the line luminosity of the nebula is predominantly determined by the ionizing luminosity of the central engine.  If the mass of cold gas increases with redshift, then the increasingly powerful AGN present in the 3CR sample at higher redshift, will illuminate a larger ionized nebula. If the velocities are predominantly gravitational, then the most powerful sources which are at the highest redshifts in our sample, and show the highest emission line gas velocity, would have to be found in very deep gravitational potentials, such as cluster environments. Alternatively, at the highest redshifts, powerful radio galaxies may be in environments which are in a relatively young evolutionary state where the kinematics are still dominated by infall. In the gravitational interpretation, the very highest velocity systems, with apparently disjoint multiple velocity systems in the nebula of a given source, are cases where two galaxies are merging or interacting, and gas systems associated with each galaxy are seen."
    },
    "0002/astro-ph0002435_arXiv.txt": {
        "abstract": "Two parameters are developed to analyze the CCD images from ground-based and/or space telescopes. The first parameter, deduced from the intensity profile of the object $sharp$, is useful to resolve stars and hot pixels. The second parameter $shape$ distinguishes the stars from the background cosmic-ray events using geometric characteristics defined by its shapes. The parameters are applied to a simulated OMC/INTEGRAL image and a HST image. ",
        "introduction": "The general problem of the fully automatic star detecting algorithms on the CCD images is the successful recognition of real star images or star-like sources in general from other types of objects and defects. These false objects originate from non-ideal characteristics of the CCD instrument (noise, bad column and hot pixels) or from the environment (nebulae, galaxy or cosmic ray events known as 'cosmics'). The frequent requirement is to find real star images on a diffuse background or in a crowded star field. Therefore, the detecting algorithm cannot use a simple intensity comparing principle, because the intensity of non-stellar objects can be comparable to star intensities.  A number of different methods for the automatic star detection on astronomical images has been used during the last two decades. For example, the method based on wavelet transforms applied to satellite ROSAT image processing was published by Damiani~et~al.~\\cite{Damiani}. The automatic  search and classification were implemented in FOCAS routine by Valdes~et~al.~\\cite{Valdes}. The PC~Vista by Trefers and Richmod~\\cite{Treffers} is a package for photometry with a simple algorithm for star detection. The well-known program DAOPHOT by Stetson~\\cite{Stetson} contains two parameters $sharp$ and $round$ and a highly robust algorithm for star detection. The determination of these parameters originates from specific characteristics of the sources on CCD images, but both parameters are not sensitive enough to resolve stars and cosmics on the CCD images exposed in focus of a space telescope and/or camera such as the OMC experiment on-board INTEGRAL (Hermsen and Winkler~\\cite{integral}). ",
        "conclusions": "The two parameters $sharp$ and $shape$ described in this paper are widely applicable to the robust detection of the star-like sources in images provided by array detectors. The primary reason for deriving the parameters was the development of a reliable automatic algorithm for the recognition of stars on CCD images. We can choose following criterion for this purpose it the intensity of the central pixel of the examined object is above an arbitrarily defined 'level of significance' (threshold) and the object lying far from the edge of the frame. Then \\begin{itemize} \\item [(a)] the object is a star for: $0 < sharp <1$ and $ 0 < shape < 1$. \\item [(b)] the object is a cosmics for: $ shape > 1$. \\item [(c)] the object is a bad column for: $ shape > 1$. \\item [(d)] the object is a hot pixel for: $ sharp > 1$. \\end{itemize} This definition provides a strict mathematical criterion for detection and classification of the various kinds of objects on CCD images.  We might specify a range of the $sharp$ and $shape$ parameters for a star centered on a pixel as $0 < sharp <1$, $ 0 < shape < 1$. The range of $sharp$ for a star centered on the edge between two pixels decrease to $0 < sharp <1/2$ and one for a star centered on the corner of four pixels decrease to $0 < sharp <1/4$, but the range for $shape$ increase to $ 0 < shape < 2$ in this both cases."
    },
    "0002/astro-ph0002090_arXiv.txt": {
        "abstract": "We describe the expected distribution of intensity for a scintillating source of finite size observed through a scattering medium, including systematic and instrumental effects. We describe measurements of the size of the Vela pulsar, using this technique. ",
        "introduction": "Waves from a pointlike source observed through a scattering medium will suffer random phase changes.  If the phase changes are much larger than 1 radian, the observer will receive radiation from many Fresnel zones, and the scattering is said to be ``strong''.  In this case the electric field at the plane of the observer is the sum of the electric field from many lines of sight, differing random phases (Goodman 1985). The net electric field is the result of a random walk.  The electric field is thus drawn from a Gaussian distribution.  Its square modulus, the intensity, is drawn from an exponential distribution (Scheuer 1968).  The region from which the observer receives radiation is known as the scattering disk.  Scattering changes phases in the Fresnel zones, and thus acts somewhat like a lens.  If the source is resolved by this ``lens'', the observed intensity is an incoherent sum from each part of the source. For a source of small but finite size, the resulting distribution of intensity is the sum of 3 exponentials.  The scales of the smaller exponentials are approximately the size of the source along either direction on the sky, in units of the linear resolution of the scattering disk (Gwinn et al. 1998).  Figure 1 shows example of the resulting intensity distributions for a point source, and for a small but resolved source.  When the source is resolved, the lowest intensities are absent.  \\begin{figure} \\plottwo{fig1a.eps}{fig1b.eps} \\caption{Left: Expected distribution of intensity for a point source in strong scintillation (dashed line); and for a source of small but finite size (solid line). From Gwinn et al. (1998). Right upper: Observed distribution of correlated flux density on a short baseline, for the Vela pulsar. Lower: Histogram shows residual to the best-fitting distribution for point source, taking into account the expected noise level. Solid curve shows best-fitting model including source size. From Gwinn et al. (2000a). } \\end{figure} ",
        "conclusions": ""
    },
    "0002/astro-ph0002059_arXiv.txt": {
        "abstract": "We report results from a program of high-resolution spectral mapping of rotational H$_2$ emission from bipolar planetary and pre-planetary nebulae. Long-slit spectra obtained with the NOAO Phoenix near-infrared spectrometer allow us to probe the molecular kinematics of these nebulae at moderate spatial resolution. We find strong evidence of a component of rotation in the equatorial H$_2$ emission from the Egg nebula (RAFGL 2688). In this nebula and in the pre-planetary nebula RAFGL 618, the H$_2$ kinematics point to the recent emergence of high-velocity polar flows, which likely mark the fairly sudden terminations of the red giant phases of their central stars. The classical bipolar planetary NGC 2346 displays distinct kinematic components, which we interpret as arising in the morphologically distinct equatorial and polar regions of the nebula. The H$_2$ rings observed in the Phoenix position-velocity maps of this nebula support the hypothesis that ring-like planetaries that display H$_2$ emission possess bipolar structure. ",
        "introduction": "The presence of molecular hydrogen emission is now recognized as a reliable indicator of bipolar structure in planetary nebulae (Zuckerman \\& Gatley 1988; Kastner et al.\\ 1994, 1996). While the polar lobes often display H$_2$, the molecular emission is, with few exceptions, brightest toward the waists of bipolar planetaries. These molecule-rich regions of planetary nebulae (PNs) appear to be the remnants of circumstellar disks or tori formed during previous, asymptotic giant branch (AGB) or post-AGB phases of the central stars. Furthermore, the available evidence suggests that the onset of H$_2$ emission postdates the AGB stage but precedes the formation of the PN (Weintraub et al.\\ 1998). This onset likely signals the beginning of a high-velocity, collimated, post-AGB wind, which shocks the previously ejected, ``slow,'' AGB wind and thereby produces the observed H$_2$ emission (Kastner et al.\\ 1999).  These observations make clear that further investigations of H$_2$ emission are important to our understanding of the origin of bipolarity in PNs. It is of particular interest to establish whether the spatially distinct waist and lobe H$_2$ emission regions are kinematically distinct as well and, furthermore, whether the kinematics bear evidence of the presence of circumstellar disks and/or high-velocity polar flows. To this end, we have undertaken a program of spectroscopic mapping of near-infrared H$_2$ emission from planetary and pre-planetary nebulae at high spectral resolution. First results from this program were presented in Weintraub et al.\\ (1998), in which H$_2$ emission was detected from a pair of bipolar pre-planetary nebulae (PPNs), and in Kastner et al.\\ (1999), where we described preliminary results for the seminal PPN RAFGL 2688. Here we present further analysis and interpretation of H$_2$ velocity mapping of RAFGL 2688, as well as H$_2$ velocity mapping results for the PPN RAFGL 618 and the bipolar planetary nebula NGC 2346 (see also Arias \\& Rosado, in this volume). ",
        "conclusions": "\\begin{figure}[htbp]  \\plotone{egg_vfield.eps}  \\caption{Comparison of the model and observed H$_2$ velocity fields of RAFGL 2688. The velocity greyscale bar applies to the center and right panels. The HST/NICMOS H$_2$ image (Sahai et al.\\ 1998) is shown at left. The observed velocity field (center) consists of velocity centroids calculated from the Phoenix data cube. In the model (right), we set the equatorial expansion velocity at $v_e=5$ km s$^{-1}$ and the equatorial rotation velocity at $v_r=10$ km s$^{-1}$. The comparison indicates that there is reasonable qualitative and quantitative agreement between model and data for this choice of parameters. }  \\end{figure}  \\subsection{RAFGL 2688}  Kastner et al.\\ (1999) presented selected velocity planes from the RAFGL 2688 Phoenix data cube. The four principal ``lobes'' of H$_2$ emission seen in direct H$_2$ images (e.g., Sahai et al.\\ 1998) are also apparent in these Phoenix data, with one pair oriented parallel to the polar axis (roughly N-S) and one perpendicular (roughly E-W).  Each of these H$_2$ lobe pairs displays a velocity gradient, with the N and E lobes blueshifted by up to $\\sim30$ km s$^{-1}$ and the S and W lobes similarly redshifted. However, the N-S and E-W H$_2$ lobe pairs differ in their detailed kinematic signatures (Kastner et al.).  The H$_2$ kinematic data for RAFGL 2688, like velocity maps obtained from radio molecular line emission, can be described in terms of a multipolar system of purely radially directed jets (Cox et al.\\ 1997; Lucas et al., these proceedings).  Given the constraints imposed by Phoenix and {\\it Hubble Space Telescope} data, however, this model would require that the ``equatorial'' components located east and west of the central star are in fact directed well above and below the equatorial plane, respectively (Kastner et al.\\ 1999).  If one postulates instead that the E-W H$_2$ emission lobes are confined to the equatorial plane of the system --- a hypothesis that appears to be dictated by certain details of the H$_2$ surface brightness distribution, as well as by simple symmetry arguments --- then one must invoke a model combining radial expansion with a component of azimuthal (rotational) velocity along the equatorial plane (Kastner et al.). In a forthcoming paper we will compare these two alternative models in more detail.  Here, we describe a specific formulation of the latter (expansion $+$ rotation) model that reproduces many of the salient features of the Phoenix data.  To construct this empirical model of the H$_2$ kinematics of RAFGL 2688, we are guided by the basic results described above. That is, the polar lobes are characterized by velocity gradients in which the fastest moving material is found closest to the star, and the slowest moving material is found at the tips of the H$_2$ emission regions. For simplicity, we assume this behavior can be described by an inverse power law relationship between velocity and radius. For the equatorial plane H$_2$ emission, meanwhile, we assume a combination of azimuthal (rotation) and radial (expansion) velocity components, whose magnitudes we denote by $v_r$ and $v_e$, respectively.  To constrain these model parameters, we compared model velocity field images with a velocity centroid image which we obtained from the Phoenix data cube. For the polar  lobes, we find that the exponent of the inverse power law velocity-distance relationship is roughly $\\sim0.7$ and that the outflow velocities at the tips of the N and S lobes are $\\sim20$ km s$^{-1}$. For the equatorial regions, good agreement between model and data is obtained for values of $v_e$ and $v_r$ that lie in the range $5-10$ km s$^{-1}$, with the additional constraint $v_e + v_r \\sim 15$ km s$^{-1}$. An example of the results for a representative model (with $v_e = 5$ km s$^{-1}$ and $v_r = 10$ km s$^{-1}$) is displayed in Fig.\\ 1. There is clear qualitative agreement between the model and observed velocity images for these parameter values, in the sense that the overall distribution of redshifted and blueshifted emission is captured by the model. Furthermore, this model reproduces specific details of the observed H$_2$ velocity distribution, such as the approximate magnitudes and positions of the velocity extrema in the four H$_2$ lobes. While this model is by no means unique, the comparison of calculated and observed velocity fields provides further support for a component of azimuthal velocity along the equatorial plane of RAFGL 2688, and offers an indication of the magnitude of this ``rotational'' component relative to the components of radial expansion both parallel and perpendicular to the polar axis of the system.  \\subsection{RAFGL 618}  \\begin{figure}[ht]  \\plotone{afgl618.eps} \\caption{Phoenix spectral image of RAFGL 618. The velocity scale of the image is centered on the systemic velocity of RAFGL 618. East is to the left. The image is displayed in a logarithmic greyscale to bring out the line wing emission, which extends to at least $\\sim \\pm 100$ km s$^{-1}$. The vertical band across the image at RA offset $\\sim0''$ is produced by continuum emission from the vicinity of the central star.}  \\end{figure}  The Phoenix spectral image obtained for RAFGL 618 is displayed in Fig.\\ 2. Bright H$_2$ emission is detected along the entire polar axis of RAFGL 618. These data demonstrate further that very high velocity H$_2$ emission is present in this bipolar outflow. The highest velocity molecular material is found closest to the central star of RAFGL 618. We conclude that the velocity gradients along the polar axes of both RAFGL 2688 and RAFGL 618 trace rapid transitions from the ``slow,'' spherically symmetric winds of their AGB progenitors to faster, collimated, post-AGB winds (Kastner et al.\\ 1999).  \\subsection{NGC 2346}  \\begin{figure}[ht]  \\plotone{ngc2346cob.eps}  \\caption{Left: Image of NGC 2346 in the 2.12 $\\mu$m line of H$_2$ obtained with the NOAO Cryogenic Optical Bench (COB; Kastner et al.\\ 1996). Right: Central region of the COB image, illustrating the slit positions used for Phoenix observations.}  \\end{figure}  \\begin{figure}[ht]  \\plotone{ngc2346phx.eps}  \\caption{Phoenix position-velocity images of NGC 2346 obtained at the slit positions illustrated in Fig.\\ 3. Top panels: images obtained as the slit was stepped northward. Bottom panels: images obtained as the slit was stepped southward. The image in the leftmost panels in each series was obtained with the slit centered on the waist of the nebula. The vertical band across the images at offsets of $0''$ and $2''$S is continuum emission from the binary companion to the central star (see Bond, these proceedings).}  \\end{figure}  In Fig.\\ 3 we display an H$_2$ image of NGC 2346 obtained with the NOAO Cryogenic Optical Bench (Kastner et  al.\\ 1996) and we illustrate the slit positions used for Phoenix spectroscopic observations. Phoenix spectral images of NGC 2346 obtained at these positions are presented in Fig.\\ 4. These images demonstrate that the H$_2$ emission from the bipolar NGC 2346 forms rings or ellipses in position-velocity space, an observation that reinforces our prior conclusion that ring-like planetaries which display H$_2$ are bipolar in structure (Kastner et  al.\\ 1994, 1996).  The position-velocity ellipse represented in the spectral image obtained with the slit at $0''$ offset (leftmost panels) is noticably tilted, with the largest redshifts found $\\sim15''$ to the east and the largest blueshifts $\\sim15''$ to the west of the central star. It is apparent from Fig.\\ 3 that this tilt is due to the orientation of the Phoenix slit with respect to the object. That is, to the east of the star the slit takes in portions of the rearward-facing (redshifted) south polar lobe of the nebula, whereas to the west the slit samples portions of the forward-facing (blueshifted) north polar lobe.  Furthermore, the position-velocity ellipses in Fig.\\ 4 contain two distinct kinematic components: a central ring associated with lower velocity material in the nebular waist and a pair of rings associated with higher velocity material in the bipolar outflow lobes. The central ring is centered at the systemic velocity of the nebula and is most apparent in the spectral images obtained near the position of the central star (i.e., in the four lefthand panels).  The southern ring is primarily redshifted (righthand bottom panels) while the northern ring (righthand top panels) is primarily blueshifted. All three rings are present in the images obtained nearest the position of the central star (lefthand panels), whereas the images obtained further from the central star display emission from only a portion of the central ring and one of the outer rings. Hence Figs.\\ 3 and 4 indicate that the H$_2$ emission from the nebula's waist produces the inner position-velocity ring, while the outer rings arise from H$_2$ emission from the polar lobes. Because of the tilt of the slit with respect to the waist of the nebula, a given slit position samples both the waist region and one or both polar lobes, resulting in a superposition of these kinematic features in a given spectral image.  In summary, the Phoenix spectral images of NGC 2346 provide strong evidence for distinct kinematic components in this nebula. These components consist of an equatorial ring or disk which is expanding at relatively low velocity ($\\sim15$ km s$^{-1}$ projected along our line of sight; Fig.\\ 4, leftmost panels) and polar lobes that are expanding at larger velocities (Fig.\\ 4, rightmost panels). Put differently, the equatorial confinement that is apparent in the morphology of this classical bipolar PN has a direct kinematic counterpart. It is tempting, therefore, to conclude that the pinched waist of NGC 2346 has its roots in processes which we are now beginning to explore in objects such as RAFGL 2688."
    },
    "0002/astro-ph0002215_arXiv.txt": {
        "abstract": "We present exactly simultaneous X-ray and TeV monitoring with {\\it RXTE} and HEGRA of the TeV blazar Mrk 501 during 15 days in 1998 June. After an initial period of very low flux at both wavelengths, the source underwent a remarkable flare in the TeV and X-ray energy bands, lasting for about six days and with a larger amplitude at TeV energies than in the X-ray band. At the peak of the TeV flare, rapid TeV flux variability on sub-hour timescales is found. Large spectral variations are observed at X-rays, with the 3--20~keV photon index of a pure power law continuum flattening from $\\Gamma=2.3$ to $\\Gamma=1.8$ on a timescale of 2--3 days. This implies that during the maximum of the TeV activity, the synchrotron peak shifted to energies $\\gtrsim 50$ keV, a behavior similar to that observed during the longer-lasting, more intense flare in 1997 April. The TeV spectrum during the flare is described by a power law with photon index $\\Gamma=1.9$ and an exponential cutoff at $\\sim$ 4~TeV; an indication for spectral softening during the flare decay is observed in the TeV hardness ratios. Our results generally support a scenario where the TeV photons are emitted via inverse Compton scattering of ambient seed photons by the same electron population responsible for the synchrotron X-rays. The simultaneous spectral energy distributions (SEDs) can be fit with a one-zone synchrotron-self Compton model assuming a substantial increase of the magnetic field and the electron energy by a factor of 3 and 10, respectively. ",
        "introduction": "BL Lacertae objects (BL Lacs) are radio-loud AGN dominated by non-thermal continuum emission from radio up to $\\gamma$-rays (MeV to TeV energies) from a relativistic jet oriented at small angles to the observer (e.g., Urry \\& Padovani 1995). While the radio through UV/X-ray continuum is almost certainly due to synchrotron emission from relativistic electrons in the jet (Ulrich, Maraschi, \\& Urry 1997 and references therein), the origin of the luminous $\\gamma$-ray radiation from BL Lacs is still uncertain. Possibilities include inverse Compton scattering of ambient photons off the jet electrons (Maraschi et al. 1992; Sikora, Begelman, \\& Rees 1994; Dermer et al. 1996), or hadronic processes (e.g.\\ Dar \\& Laor 1997; Mannheim 1993).  A breakthrough was provided by the discovery of TeV emission from a handful of such sources, all characterized by a synchrotron peak at higher energies (High-energy peaked BL Lacs, or HBLs). One of these is Mrk 501 ($z$=0.034). This source came into much attention after it exhibited a prolonged period of intense TeV activity in 1997 (Catanese et al. 1997; Hayashida et al.\\ 1998; Quinn et al.\\ 1999; Aharonian et al. 1997,1999a-c; Djannati-Atai et al. 1999), accompanied by correlated X-ray emission on timescales of days. Interestingly, this exceptional TeV activity was accompanied by unusually hard X-ray emission up to $\\gtrsim$ 100 keV (Pian et al. 1998a; Catanese et al. 1997; Lamer \\& Wagner 1999; Krawczynski et al.\\ 1999), unprecedented in this or any other BL Lac. The hard X-ray spectrum implied a shift toward higher energies of the synchrotron peak, usually located at UV/soft X-rays (e.g., Sambruna, Maraschi, \\& Urry 1996; Kataoka et al. 1999), by more than three decades, persistent over a timescale of $\\sim$ 10 days (Pian et al. 1998a). Further observations with {\\it BeppoSAX} in April-May 1998 and in May 1999 during periods of TeV lower flux showed that the synchrotron peak had decreased to $\\sim$ 20 and 0.5 keV, respectively (Pian et al. 1998b, 1999). These secular variations of the synchrotron peak suggest a powerful mechanism of particle energization, operating over timescales of years.  Because of its bright TeV emission and unusual X-ray spectral properties, we selected Mrk 501 for an intensive monitoring in 1998 June using HEGRA and the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}), with a sampling designed to probe correlated variability at the two wavelengths on timescales of one day or shorter. Here we report the first results of the campaign, which is characterized by the detection of a strong flare at both TeV and X-ray energies after a period of very low activity. The structure of this paper is as follows.  We describe the sampling and the observations in \\S~2, the X-ray and TeV light curves in \\S~3.1, and the TeV and X-ray spectra in \\S\\S~3.2--3.3.  Implications of the data are discussed in \\S~4. ",
        "conclusions": "Since the typical flux variability timescale of Mrk 501 in TeV $\\gamma$-rays and X-rays can be much less than one day, it is important to have truly simultaneous observations in both bands. It is also important to have reasonably continuous sampling on timescales of at least one day in order to have an accurate picture of the dynamics of the source.  For this reason, we conducted a 15-day TeV/X-ray monitoring with diurnal {\\it RXTE} observations exactly in the HEGRA visibility windows.  After 10 days of quiescence, the source exhibited a strong flare at both TeV and X-rays lasting six days, with a flux exceeding the pre-flare level by a factor of $\\sim $~20 at TeV energies during a 2-day maximum, and with lower amplitudes (factor 2--4) at X-rays.  We also report the first detection of TeV flux variability on sub-hour timescales in Mrk 501 (\\S~3.1).  By chance, our multiwavelength campaign in 1998 June coincided with the only high TeV activity of the source during that year. Luckily, we were able to follow the evolution of the TeV flare not only during the pre-flare and flare stage but also during the decay stage. The TeV spectrum during the flare is similar to the spectra observed in 1997, suggesting that the flaring episode we witnessed in 1998 June was a scaled-down version of the longer-lasting 1997 flare.  This conclusion is bolstered by the strong spectral variations we observe in the X-rays. Our {\\it RXTE} observations show that the X-ray continuum in 3--20 keV flattened by $\\Delta\\Gamma_{3-20~keV} \\sim 0.5$ from the beginning of the campaign ($\\Gamma_{3-20~keV}=2.3$) to the flare maximum ($\\Gamma_{3-20~keV}=1.8$).  Interestingly, at the peak of the TeV flare the X-ray slope was similar to the 2--10 keV slope measured in 1997 April, May, and July with {\\it BeppoSAX} and {\\it RXTE} (Pian et al. 1998a; Lamer \\& Wagner 1999; Krawczynski et al. 1999).  This implies a similar shift of the synchrotron peak frequency at higher energies, $\\gtrsim 50$ keV (Figure 3).  While in April 1997 the X-ray continuum flattened by 0.4 within approximately two weeks, we see here a comparable flattening within only $\\sim$ 2--3 days. Note that large changes of the position of the synchrotron peak are relatively rare. Besides Mrk 501, they were observed to-date only in two HBLs, 1ES 2344+514 (Giommi, Padovani, \\& Perlman 1999) and 1ES 1426+428 (Ghisellini, Tagliaferri, \\& Giommi 1999), but not in Mrk 421, PKS 2155--304, or any other BL Lac. Our observations provide the first evidence that in Mrk 501 the synchrotron peak may change on relatively short timescales ($\\sim$ a few days).  Several models have been suggested to explain the TeV radiation from blazars. A popular scenario are the leptonic models, where TeV $\\gamma$-rays are produced via inverse Compton scattering of directly accelerated electrons on external and/or internal photons (e.g., Sikora 1997).  For Mrk 501, an object without strong broad line emission, the synchrotron self-Compton (SSC) model is almost commonly accepted as the most probable explanation for the observed X-ray/TeV-$\\gamma$-ray emission (e.g., Tavecchio et al. 1998; Kataoka et al. 1999).  Presently, the SSC model is the only model (at least in its simplified, ``one-zone'' version) which has been developed to a level which allows conclusive predictions which can be compared with experimental results.  In particular, the SSC scenario is able to give satisfactory fits to both the X-ray and the TeV spectra (Pian et al. 1998a; Hillas 1999; Krawczynski et al. 1999). We used the code developed by Coppi (1992) to fit our simultaneous SEDs in Figure 3, assuming emission from a one-zone, homogeneous region and incorporating Klein-Nishina effects. The key parameters used in this model are the Doppler factor $\\delta_{\\rm j}$ of the relativistic plasma, the radius $R$ of the emission region, the magnetic field $B$, and the electrons' maximum energy $E_{max}$.  The results of the fits are shown in Figure 3 as solid lines, and the parameters' values are reported in the caption.  As discussed further below, the models were computed without correcting for the extragalactic extinction of TeV photons due to $\\gamma$/$\\gamma$ pair production with the photons of the Diffuse Extragalactic Background Radiation (e.g.\\ Gould \\& Schr\\'{e}der 1966).  In the lower panel, we plot the ratio of the data and best-fit model between the high-state and pre-flare. The latter plot emphasizes that, while the TeV spectra of both states are quite similar, the X-ray spectra of the pre-flare and the flare state are significantly different.  In SSC models the hardening of the X-ray spectrum during the flare can be attributed to a shift of the peak frequency $\\nu_{\\rm s}$ of the synchrotron radiation, $\\nu_{\\rm s}\\propto~B \\times E^2_{\\rm max}$.  Assuming an increase of both the magnetic field and the maximum energy during the flare, the dramatic changes of the X-ray spectrum are readily explained (see Figure 3).  While the increase of magnetic field does not affect the $\\gamma$-ray spectrum, the increase of $E_{\\rm max}$ does make the inverse Compton (IC) spectrum harder.  However, since the $\\gamma$-rays are produced in the Klein-Nishina regime, this effect is less pronounced in IC than in the synchrotron radiation component.  The fits shown in Figure 3 correspond to the following model parameters: $B$=0.03~G, $E_{\\rm max}=2$~TeV (exponential cutoff energy), $R=4 \\times 10^{16}$\\,cm for pre-flare state, and $B$=0.1~G, $E_{\\rm max}=20$~TeV, $R=2.7 \\times 10^{15}$\\,cm for the flare state. For both cases a Doppler factor of $\\delta_{\\rm j}=25$ is assumed. Note that the latter value of the Doppler factor implies that internal absorption of the TeV $\\gamma$-rays by lower frequency photons can be completely neglected (e.g.\\ Celotti, Fabian, \\& Rees 1998). Furthermore, the chosen Doppler factor and radius of the emitting volume in the flaring state imply time variability down to $t=R/(c~\\delta_{\\rm j}) =$1~hour which agrees with the observed flux variability following from Figure 2.  For the flare state with good statistics up to $\\sim$ 10~TeV, the model over-predicts the TeV flux above $\\sim$ 5~TeV, in particular by a factor of $\\sim$ 2.5 at 10~TeV. This discrepancy should not be overemphasized, but could well be the result of intergalactic extinction due to $\\gamma$/$\\gamma$ pair production.  In summary, we have performed a 2-week monitoring campaign of the HBL Mrk 501 in 1998 June with HEGRA and {\\it RXTE}, with a sampling designed to probe TeV/X-ray correlation on timescales of several hours. We detected a strong flare at both wavelengths, rising from a period of very low activity, well correlated at TeV and X-rays on time scales of $\\lesssim$ 1 day, accompanied by large ($\\Delta\\Gamma_{3-20~keV} \\sim 0.5$) spectral variability at X-rays. Our results support an interpretation in terms of a canonical synchrotron-self Compton scenario. Future campaigns with a more intensive sampling designed to probe correlation on shorter time scales at both X-ray and TeV energies are needed to set more stringent constraints on the radiative processes which play an important role in the evolution of the flare."
    },
    "0002/astro-ph0002509_arXiv.txt": {
        "abstract": "We consider some aspects of the evolution of massive stars which can only be elucidated by means of ``indirect\" observations, \\ie measurements of the effects of massive stars on their environments. We discuss in detail the early evolution of massive stars formed in high metallicity regions as inferred from studies of HII regions in external galaxies. ",
        "introduction": "Massive stars play a crucial role in the evolution of galaxies and the whole Universe, because they are the primary sources of radiative ionization and heating of the diffuse medium, they provide most of the nucleosynthesis products to boost the metal content of galaxies and the intergalactic medium, and they constitute a major supply of kinetic energy for galaxies, both through stellar winds during their quiescent phases and, eventually, in the form of fast ejecta from supernova explosions.  Therefore, it is fundamental to reach a proper understanding of the formation processes, the detailed properties and the evolution of massive stars.  Despite the fact that their high luminosities make them ``easy\" targets for detailed observational studies, many aspects and properties of massive star evolution are far from being fully understood. This is because in any stellar generation, massive stars constitute a small fraction of the newly formed stars (say, less than 1\\% by number), they are ``elusive\" in that their lifetimes are very short (say, less than 10--20 million years), and often they are heavily obscured by the parent molecular clouds where they were formed, making even their identification rather cumbersome.  As a consequence it is not easy to cover all evolutionary phases with {\\it direct} observations of a statistically significant sample of objects.  One can overcome these difficulties and gain additional insights by considering phenomena that {\\it indirectly} can provide hints and clues to the problem. In other words, besides studying individual massive stars, one can look at the effects that these stars have on their environments (\\eg HII regions, circumstellar nebulae, SNRs), and infer from there what the stars were doing in special phases of their evolution (\\eg formation, LBV and pre-SN phases, etc.) that would not be accessible in other ways.  Thus, one can use radio observations of supernovae, which probe the circumstellar material ejected by the progenitor stars several thousand years before explosion, to study the very last phases of their evolution. These phases represent a tiny fraction of a massive star lifetime, $\\sim$0.1\\%, and, therefore, they are extremely difficult to reveal and study with direct observations.   Although this is an interesting aspect, we are not going to review it here, but rather we refer the reader to recent papers (Montes \\etal~ 1998, Panagia \\etal~ 1999, Weiler \\etal~ 1999).  Here, we consider and discuss one  particular aspect of the evolution of massive stars, namely their formation and early evolution in high metallicity environments.  We will show that observations of HII regions in external galaxies show that the ionization of He is much lower than that of H when the O/H ratio in the gas is appreciably higher than solar. This implies that at high metallicities either very massive stars ($M>25~M_\\odot$) do not form, or they never reach their expected ZAMS location. ",
        "conclusions": "Ionized helium abundances are shown in Figure 1 as a function of oxygen abundances, for all of the H~II regions in spiral and irregular galaxies and the Blue Compact Galaxies of our sample. The long baseline in metallicity offers a unique opportunity to constrain {\\it both } the abundance of primordial helium Y$_P$ {\\it and} to the \\dYdZ gradient, thus fully determining the ``helium enrichment curve'' (HEC) which relates the total abundance of helium to the abundance of oxygen. Considering that evolutionary effects always decrease the  He$^+$/H$^+$ ratio because the aging of a stellar cluster results in a softening of the ionizing radiation field,  and that young clusters are observationally favoured because they are intrinsically brighter than older clusters, the HEC can be derived by determining the upper envelope of the distribution shown in Figure 1. Among the class of curves $Y = Y_0 + \\Delta Y/\\Delta Z \\times Z$ the best fit to the upper envelope of the observations is obtained for $Y_0$ = 0.243 and (\\dYdZ)$_{\\odot}$ = 3.2 (see dashed curve in Fig.~1). This relation is consistent with the observational results of Pagel \\etal~ (1992) as well as with Maeder's (1992) theoretical models.  An inspection to Figure 1 reveals that the observed He$^+$/H$^+$ ratio appears to be almost constant up to solar O abundance (log(O/H)$_\\odot+12\\simeq$8.8) and then it declines rather quickly for higher metallicities.  This is a clear sign that He is progessively less ionized as the O abundance increases, and implies that the mean radiation temperature of the ionizing stars becomes lower than about 38,000~K around log(O/H)$\\simeq8.5$.  \\begin{figure} \\centerline{ \\psfig{figure=fig_vulchem.ps,width=9cm,height=9cm} } \\caption{ He$^+$/H$^+$ ratios (by number) as a function of the O~abundances for the 287 H~II regions in 46 Spiral and Irregular galaxies (filled symbols) and for the 87 Blue Compact Galaxies (open squares) of our sample. The cross in the lower left corner represents the typical uncertainty of the data. The boxes correspond the $\\pm1\\sigma$ deviations from the averages calculated within 0.2 dex intervals in log(O/H). The dashed line represents the helium enrichment curve (HEC) and the solid line is a model calculation in which the He ionization decreases exponentially with increasing O~abundance.} \\end{figure}  We find that mechanisms through which metallicity acts ``directly'' on the radiation field of the ionizing star cluster are not enough to explain the observed gradient of the He ionization fraction with O abundance. Most of the effect appears instead to be due to ``indirect'' mechanisms, \\ie a marked deficiency of hot stars with increasing metal abundances. There are at least three possible scenarios to explain this fact:\\\\ 1. The IMF slope becomes steeper for higher metallicities.\\\\ 2. The IMF upper cutoff moves to lower masses for higher metallicities.\\\\ 3. The most massive stars become progressively unable to provide ionizing radiation, either because at high metallicities the remnant of their pre-MS cocoons remains optically thick over most of a star's lifetime, or because pulsational instabilities prevent the most massive stars from reaching their expected ZAMS surface conditions. \\\\ Our model calculations show that varying only the slope of the IMF, \\ie point (1), does not give a satisfactory fit to the data because the resulting ionization decline would be too shallow. On the other hand, point (2), \\ie a systematic variation of the IMF upper mass cut-off with metal abundance ($m_U \\propto$ Z$^{-\\beta}$, $\\beta > 0$) can reproduce the observed trend of the He ionization, with $m_{U\\odot}$ = 48 \\Msun and $\\beta$ = 0.60. Point (3) could also account for the observations provided that the invoked effects are indeed capable to produce the sharp decline of He ionization as observed.  From the observational point of view there are no direct studies to conclusively discriminate between points (2) and (3). Observations of massive stars near the Galactic Center, such as the Pistol star, the Sickle and the Quintuplet clusters (\\eg Figer \\etal~ 1999 and references therein) seem to favor the third possibility, because they are so bright ($log(L/L_\\odot)>6$) that they must be quite massive. On the other hand, one may argue that those clusters are so close to the Galactic Center that tidal forces may drastically affect the dynamical processes that lead to the formation of stars and, therefore, they may not be representative of normal situations. The ideal investigation to clarify this issue should include nebular and stellar spectroscopy of a large sample of HII regions in galaxies which display marked effects  of incomplete He ionization, such as M51 or M83.  Another discriminant between hypotheses (2) and (3) is that if a lowering of the IMF upper cutoff is {\\it the } explanation (\\ie point (2)), then the frequency of Wolf-Rayet stars relative to early type stars is expected to be abnormally low in the high-Z regions because in this case the reduced ionization is entirely due to the lack of truly massive stars that are expected to end up as WR stars in their final stages of evolution."
    },
    "0002/astro-ph0002023_arXiv.txt": {
        "abstract": "We present simultaneous ASCA and RXTE data on the bright Seyfert 1 galaxy IC4329a. The iron line is significantly broadened, but not to the extent expected from an accretion disk which extends down to the last stable orbit around a black hole. We marginally detect a narrow line component, presumably from the molecular torus, but even including this gives a line profile from the accretion disk which is significantly narrower that that seen in MCG--6--30--15, and is much more like that seen from the low/hard state galactic black hole candidates.  This is consistent with the inner disk being truncated before the last stable orbit, forming a hot flow at small radii as in the ADAF models. However, we cannot rule out the presence of an inner disk which does not contribute to the reflected spectrum, either because of extreme ionisation suppressing the characteristic atomic features of the reflected spectrum or because the X--ray source is intrinsically anisotropic, so it does not illuminate the inner disk.  The source was monitored by RXTE every 2 days for 2 months, and these snapshot spectra show that there is intrinsic spectral variability. The data are good enough to disentangle the power law from the reflected continuum and we see that the power law softens as the source brightens.  The lack of a corresponding increase in the observed reflected spectrum implies that either the changes in disk inner radial extent/ionisation structure are small, or that the variability is actually driven by changes in the seed photons which are decoupled from the hard X--ray mechanism. ",
        "introduction": "Accretion of material onto a black hole is known to produce hard X--ray ($E \\ge 2$ keV) emission: satellite observations over the last 30 years have conclusively shown this to be the case for both stellar mass black holes in our own Galaxy, and the supermassive black holes which power the Active Galactic Nuclei (AGN).  However, it is still the case that the mechanism by which the gravitational potential energy is converted into high energy radiation is not understood.  Standard models of an accretion disk (Shakura \\& Sunyaev 1973, hereafter SS) produce copious UV and even soft X--ray radiation, but are completely unable to explain the observed higher energy X--ray emission which extends from the disk spectrum out to 200 keV. Clearly something other than the standard model is required, and the two currently favored candidates are either magnetic reconnection above an accretion disk (e.g. Galeev, Rosner \\& Vaiana 1979; Haardt, Maraschi \\& Ghisellini 1994; di Matteo 1998), or that some part of the accretion flow is not given by the standard disk configuration.  The recent rediscovery of another stable solution of the accretion flow equations lent plausibility to this second possibility. The standard SS accretion disk model derives the accretion flow structure in the limit where the gravitational energy released is radiated locally in an optically thick, geometrically thin disk. This contrasts with the new accretion models, where below some critical mass accretion rate, $\\dm\\le \\mdcrit$, the material is not dense enough to thermalise and locally radiate all the gravitational potential energy which is released. Instead the energy can be carried along with the flow (advected), eventually disappearing into the black hole. These solutions give an X--ray hot, optically thin, quasi--spherical flow (Narayan \\& Yi 1995; Esin, McClintock \\& Narayan 1997).  Clearly it would be nice to know which (if any!) of these models for the hard X--ray emission is correct. In its most general form the problem comes down to understanding the geometry of the cool, optically thick accretion flow. If this extends down to the last stable orbit around the black hole then it is unlikely that the advective flow can exist (e.g.\\ Janiuk, \\.{Z}ycki \\& Czerny 1999). Conversely, if the optically thick disk truncates before this point then a composite model with an outer disk and inner, hot, advective flow (ADAF) may be favored.  The accretion flow can be tracked via X--ray spectroscopy. Wherever hard X--rays illuminate optically thick material then this gives rise to a Compton reflected continuum and associated iron K$\\alpha$ fluorescence line (e.g.\\ Lightman \\& White 1988; George \\& Fabian 1991; Matt, Perola \\& Piro 1991).  The amplitude of the line and reflected continuum is dependent on the solid angle subtended by the disk to the X--ray source (as well as on elemental abundances, inclination and ionisation). The two models outlined above can then be distinguished by the amount of reflection and line, since an untruncated disk should subtend a rather larger solid angle than a truncated one.  A survey of AGN showed that Seyfert 1 spectra are consistent with a power law X--ray spectrum illuminating an optically thick, (nearly) neutral disk, which subtends a solid angle of $\\sim 2\\pi$ (Pounds et al. 1990; Nandra \\& Pounds 1994). This then seems to favor the magnetic reconnection picture. However, this contrasts with the situation in the Galactic Black Hole Candidates (GBHC). These are also thought to be powered by accretion via a disk onto a black hole, and, in their low/hard state, show broad band spectra which are rather similar to those from AGN, but have an apparently rather smaller amount of reflection (e.g. Gierli\\'{n}ski et al. 1997; Done \\& \\.{Z}ycki 1999). One potential drawback of this comparison is that AGN inhabit a more complex environment than the GBHC. Unification schemes for Seyfert galaxies propose that there is a molecular torus which enshrouds the nucleus. The molecular torus can also contribute to the reflected spectrum, perhaps distorting our view of the very innermost regions in AGN.  These two potential sites for the reflected component can be distinguished spectrally: any features from the accretion disk should be strongly smeared by the combination of special and general relativistic effects expected from the high orbital velocities in the vicinity of a black hole (Fabian et al.  1989), whereas the molecular torus is at much larger distances so its reflected features should be {\\it narrow}. The seminal ASCA observation of the AGN MCG--6--30--15 showed that the line is so broad that it requires that the accretion disk extends down to {\\it at least} the last stable orbit in a Schwarzschild metric, with {\\it no} narrow component from the molecular torus (Tanaka et al. 1995; Iwasawa et al. 1996), and that the relativistically smeared material subtended a solid angle of $\\sim 2\\pi$ with respect to the X--ray source.  This very clear cut result then seems to rule out the advective flows, at least in their simplest form.  However, the GBHC again show a rather different picture: they have a line which is detectably broad, but not so broad as might be expected for a disk extending all the way down to 3 Schwarzschild radii ($R_{\\rm Schw} = 2 G M/c^2$; \\.{Z}ycki, Done, \\& Smith 1997, 1998, 1999; Done \\& \\.{Z}ycki 1999).  This {\\it is} consistent with the truncated disk geometry, and so perhaps with the advective flow models.  Only in the soft/high state do the GBHC seem to show the extreme relativistic smearing and large amount of reflection seen in the MCG--6--30--15 spectra (\\.{Z}ycki et al. 1998, Gierli\\'{n}ski et al. 1999).  Is there a real difference in accretion geometry between GBHC and AGN, pointing to a difference in radiation mechanism ? This seems unlikely, since both classes are ultimately accreting black holes. Are subtle ionisation effects masking the derived disk parameters in the GBHC (Ross, Fabian \\& Young 1999).  Or is MCG--6--30--15 unusual among AGN (and GBHC) in having such a relativistic disk?  Perhaps MCG--6--30--15 is an AGN in a state which corresponds to the soft state GBHC?  One factor supporting the latter is a recent study by Zdziarski, Lubi\\'{n}ski, \\& Smith (1999) which showed that there is a correlation between the intrinsic spectral slope, $\\Gamma$, and the solid angle subtended by the reflecting material, $\\Omega/2\\pi$. In their plots, MCG--6--30--15 is the AGN with the steepest spectrum, and highest amount of reflection. This correlation also holds for individual objects (such as NGC 5548 for AGN and Nova Muscae for the GBHC), where the intrinsic spectrum hardens as the amount of reflection decreases.  This suggests that there is a universal physical mechanism and/or geometry for both classes, with perhaps a single parameter determining the state of a given source through a feedback between the geometry and physical conditions in the X--ray emitting region. Such a feedback could be provided by e.g.\\ soft photons from the thermalized fraction of the hard X--rays intercepted by the reprocessing medium, and the parameter could be the inner disk radius.  Perhaps for MCG--6--30--15 (and other soft AGN and the soft state GBHC) the cool accretion disk extends down to the innermost stable orbit around the black hole with the X--rays being powered by magnetic reconnection above the disk, while for harder spectra AGN (and the low state GBHC) the inner disk recedes outwards, being replaced by an X--ray hot (advective ?) flow.  As the disk recedes it subtends a smaller solid angle, so there is less reflection (and less relativistic smearing), but there are also fewer seed photons from the disk (both from intrinsic emission and by reprocessing) for Compton scattering into the intrinsic power law, giving a harder intrinsic power law (\\.{Z}ycki et al. 1999;  Zdziarski et al. 1999).  If this is true, then this clearly predicts that the Fe K$\\alpha$ line {\\it in AGN} is not always as broad as in the extreme case of MCG-6-30-15 (Tanaka et al.\\ 1995).  Previous ASCA studies on a sample of objects (Nandra et al.\\ 1997) hint towards such a possibility, since a whole range of geometries were inferred for various objects.  Here we test this idea using ASCA, XTE and OSSE data from IC 4329a, the brightest 'typical' Seyfert 1 in the X--ray band (cf. Madejski et al. 1995). It lies towards the middle of the $\\Gamma-\\Omega/2\\pi$ plot of Zdziarski et al. (1999), and (consequently) has a spectrum very close to that of the mean Seyfert 1 spectrum compiled by Zdziarski et al. (1995). IC 4329a may then be used as a template for Seyfert galaxies as a class, unlike MCG--6--30--15 which has a rather steep X--ray spectral index. ",
        "conclusions": "$\\bullet$ Not all AGN have the extreme relativistic line profiles expected from a disk extending down to the innermost stable orbit around a black hole. This is consistent with either the inner disk being truncated before the last stable orbit, or with an inner disk which produces no significant reflected features either through anisotropic illumination or extreme ionisation. Simple photo--ionisation models, where the ionisation varies smoothly as a function of radius can be ruled out by the data, but these may differ substantially from more detailed models of the ionisation structure.  $\\bullet$ There is intrinsic spectral variability, where the power law softens as the source brightens. This implies that the soft seed photons are increasing faster than the increase of the hard X--ray luminosity.  The lack of a corresponding increase in the observed reflected spectrum implies that either the changes in disk inner radial extent/ionisation structure are small, or that the variability is actually driven by changes in the seed photons which are decoupled from the hard X--ray mechanism."
    },
    "0002/astro-ph0002353_arXiv.txt": {
        "abstract": "The application to main-sequence stars of the rotation-in-duced mixing theory in the presence of $\\displaystyle \\mu $-gradients leads to partial mixing in the lithium destruction region, not visible in the atmosphere. The induced lithium depletion becomes visible in the sub-giant phase as soon as the convective zone deepens enough. This may explain why the observed `` lithium dilution \" is smoother and the final dilution factor larger than obtained in standard models, while the lithium abundance variations are very small on the main sequence. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002165_arXiv.txt": {
        "abstract": "The origin of GRBs have been a mystery for almost 30 years. The afterglow observed in the last few years enabled redshift determination for a handful of bursts, and the cosmological origin is now firmly established. Though the distance scale is settled, there still remains orders of magnitude uncertainty in their rate and in the total energy that is released in the explosion due to the possibility that the emission is not spherical but jet-like. Contrary to the GRB itself, the afterglow can be measured up to months and even years after the burst, and it can provide crucial information on the geometry of the ejecta. We review the theory of afterglow from jets and discuss the evidence that at least some of the bursts are not spherical. We discuss the prospects of polarization measurements, and show that this is a powerful tool in constraining the geometry of the explosion. ",
        "introduction": "The study of $\\gamma$-ray bursts was revolutionized when the Italian-Dutch satellite BeppoSAX delivered arcminutes positioning of some GRBs, within a few hours after the event. This enabled other ground and space instruments to monitor the relatively narrow error boxes. Emission in X-ray, infrared, optical and radio, so called ``afterglow'', was observed by now for more than a dozen of bursts.  The current understanding of the GRBs phenomenon is that a compact source emits relativistic flow with Lorentz factor $\\gamma$ of at least a few hundreds. This flow emits, probably by internal shocks (see e.g. \\cite{SP97,FMN96}), the GRB. After these internal shocks have produced the GRB, the ultra relativistic flow interacts with the surrounding medium and decelerates. Synchrotron radiation is emitted by the heated surrounding matter. As more and more of the surrounding mass is accumulating, the flow decelerates and the emission shifts to lower and lower frequencies. Excitingly, the afterglow theory is relatively simple. It deals with the emission on timescales much longer than those of the GRBs. The details of the complex initial conditions are therefore forgotten and the evolution depends only on a small number of parameters.  We begin by clarifying some of the confusing terminology. There are two distinct, but related, effects.  The first, {\\bf ``jets''}, describes scenarios in which the relativistic flow emitted from the source is not isotropic but collimated towards a finite solid angle. The term jet refers to the geometrical shape of the relativistic flow emitted from the inner engine. The second effect is that of {\\bf ``relativistic beaming''}. The radiation from any object that is radiating isotropically in its own rest frame, but moving with a large Lorentz factor $\\gamma$ in the observer frame, is collimated into a small angle $1/\\gamma$ around its direction of motion. This is an effect of special relativity. It has nothing to do with the ejecta's geometry (spherical or jet) but only with the fact that the ejecta is moving relativisticly.  The effect of relativistic beaming allows an observer to see only a small angular extent, of size $1/\\gamma$ centered around the line of sight. Unfortunately, the term beaming was also used for ``jets'' by many authors (including myself). We will keep a clear distinction between the two in this paper.  Since we know the flow is ultra-relativistic (initially $\\gamma>100$), there is no question that the relativistic beaming is always relevant for GRBs. The question we are interested in is that of the existence of ``jets''.  The idealized description of a jet is a flow that occupies only a conical volume with half opening angle $\\theta_0$. In fact the relativistic dynamics is such that the width of the matter in the direction of its propagation is much smaller than its distance from the source by a factor of $1/\\gamma^2$. The flow, therefore, does not fill the whole cone. Instead it occupies only a thin disk at its base, looking more like a flying pancake \\cite{P99} - see figure \\ref{grbjets}.  If the ``inner engine'' emits two such jets in opposite directions then the total solid angle towards which the flow is emitted is $\\Omega=2\\pi \\theta_0^2$. The question whether the relativistic flow is in the form of a jet or a sphere has three important implications.  \\begin{figure*} \\centerline{\\epsfig{file=f3.eps,width=3.0in}} \\vspace{10pt} \\caption{Schematic geometric description of jets in GRBs. The scheme shows the multiple shells before internal shocks have occurred. After that they all merge to one shells with typical width a factor of $\\gamma^2$ thinner than their distance from the source. } \\label{grbjets} \\end{figure*}  \\noindent {\\bf The Total Emitted Energy.} Optical observations of afterglows enabled redshift determination, and therefore a reasonably accurate estimate of the distance, $D$, to these events (the uncertainty is now in the cosmological parameters of the universe). The so called ``isotropic energy'' can then be inferred from the fluence $F$ (the total observed energy per unit area at earth) as $E_{iso}=4\\pi D^2 F$ (taking cosmological corrections into account, $D=D_L/\\sqrt{1+z}$ where $D_L$ is the luminosity distance and $z$ is the redshift). The numbers obtained in this way range from $10^{51}$erg to $10^{54}$erg with the record of $3\\times 10^{54}$erg held by the famous GRB~990123. These huge numbers approach the equivalent energy of a solar mass, all emitted in a few tens of seconds!  These calculations assumed that the source emitted the same amount of energy towards all directions. If instead the emission is confined to some solid angle $\\Omega$ then the true energy is $E=\\Omega D^2 F$. As we show later $\\Omega$ is very weakly constrained by the GRB itself and can be as low as $10^{-6}$. If so the true energy in each burst $E \\ll E_{iso}$. We will show later that interpretation of the multi-wavelength afterglow lightcurves indeed indicates that some bursts are jets with solid angles considerably less than $4 \\pi$. The isotropic energy estimates may be fooling us by a few orders of magnitudes! Clearly this is of fundamental importance when considering models for the sources of GRBs.  \\noindent {\\bf The Event Rate.} BATSE sees about one burst per day. With a few redshifts measured this translates to about $10^{-7}$ bursts per year per galaxy. However, if the emission is collimated to $\\Omega \\ll 4\\pi$ then we do not see most of the events. The true event rate is then larger than that measured by BATSE by a factor of $4\\pi /\\Omega$. Again this is of fundamental importance. Clearly, the corrected GRB event rate must not exceeds that of compact binary mergers or the birth rate of massive stars if these are to produce the majority of the observed GRBs.  \\noindent {\\bf The Physical Ejection Mechanism.} Clearly, different physical models are needed to explain collimated and isotropic emission.  For example, in the collapsar model (e.g. \\cite{MW99}), relativistic ejecta that is capable of producing a GRB is produced only around the rotation axis of the collapsing star with half opening angle of about $\\theta_0 \\cong 0.1$. Such models would have difficulties to explain isotropic bursts as well as very narrow jets.  With these uncertainties we are therefore left with huge ignorance in how, how much and how many GRBs are produces. The question as to whether the emission of GRBs is spherical or collimated in jets is fundamental to almost all aspects of the GRB phenomenon. ",
        "conclusions": "Now when redshifts for GRBs are routinely measured, the largest uncertainty in their energy budget and event rate is the possibility that the emission is not spherical but jet-like. We discussed the theory of afterglow from jet-like event. These should produce a substantial break at all frequencies. The time where this break occurs is an indication of the jets opening angle. GRB~990510 seems to be a perfect example for this behavior.  The inferred opening angle is about $0.1$ consistent with upper limits from searches of orphan X-ray afterglows \\cite{GHV+99}. Several other candidate for jets are bursts with fast decline, where the break presumably took place before the earliest observation. This question will be settled when more frequent early observations are available. We have shown that afterglow from jets should show a unique signature of polarization, at detectable levels. Observing such signature will confirm the jet interpretation and the synchrotron model in general.  {\\bf Acknowledgements} I thank Titus Galama for very useful comments, and the Sherman Fairchild foundation for support."
    },
    "0002/astro-ph0002486_arXiv.txt": {
        "abstract": "We explore one of plausible causes of asymmetry in the wind of an AGB star -- the presence of a binary companion. We have developed a simple method for estimating the intrinsic non-sphericity of the outflow in that case. Assuming the Roche model, local stellar parameters ($T_{\\rm eff}$, $g$) are calculated, and then a formula for the local mass loss rate is applied. As a by-product a relation between the ratio of stellar radius to the Roche lobe radius and the overall wind enhancement is obtained. ",
        "introduction": "Shapes of Planetary Nebulae, that deviate from spherical symmetry (in particular -- axisymmetric ones) are often ascribed to binary interactions (e.g. Soker 1997). When an Asymptotic Giant Branch star loses its mass -- which is to become a PN -- the companion can affect the trajectories of the outflowing matter, concentrating it to the orbital plane. A density gradient between equatorial and polar regions is created, and when the star leaves AGB and the hot fast wind starts to break its way through the remnants of the expelled giant's envelope, the elliptical or bipolar symmetry forms naturally.  In attempts to calculate such effects it is widely assumed, that the intrinsic AGB wind is spherically symmetric and the asymmetry is introduced only by the companion's influence. But this need not to be the case for relatively close binaries, where the giant is noticeably distorted by tidal forces. Differences in local conditions through the stellar surface, mainly in temperature and gravity, may lead to different intensities of the outflow. Thus the wind would show an intrinsic directivity.  We investigate this possibility, using a simple model. ",
        "conclusions": "As one could expect, we find significant differences in the local mass loss rates across the surface of the giant distorted by the presence of a companion. The gradient between the equatorial and polar regions is evident, although the strongest enhancement occurs toward the companion. It would be interesting to use these results as an input for modelling shapes of the Planetary Nebulae formed in binary systems, replacing the assumption of intrinsic sphericity of the AGB wind.  For evolutionary calculations considering binary evolution it is important to know by how much the stellar mass loss rate is affected by the presence of the companion. Tout and Eggleton (1988) proposed a formula, according to which the tidal torque would enhance the mass loss by a factor of $1 + B \\times (R/R_{RL})^{6}$, where $B$ is a free parameter to be adjusted (ranging from $5 \\times 10^{2}$ to $10^{4}$).  In our model the mass loss rate enhancement depends mainly on the $g$ value at the giant's point closest to the binary companion. Therefore one may try a simple analytical approach to derive a similar relation for wind enhancement. Let us denote the gravity of a~single star by $g_{single}$. Expanding the ratio $g^{-1}/g_{single}^{-1}$ at the point closest to the companion into a series in small $R/R_{RL}$ gives $const~\\times~(R/R_{RL})^3$ as a first non-zero term following the unity. This leads to the following dependence for the mass loss rate: \\begin{displaymath} \\dot M = \\dot M_{single} \\, (1+ const \\times (R/R_{RL})^3). \\end{displaymath}  Our numerical results confirm the above relation up to $\\log R/R_{RL} \\sim -0.2$ (i.e.~$R/R_{RL} \\sim 2/3$), which is shown on Fig.~2."
    },
    "0002/astro-ph0002492_arXiv.txt": {
        "abstract": "High energy gamma ray astronomy is now a well established field and several sources have been discovered in the region from a few GeV up to several TeV. If sources involving hadronic processes exist, the production of photons would be accompanied by neutrinos too. Other possible neutrino sources could be related to the annihilation of WIMPs at the center of galaxies with black holes.  We present the results of a search for point-like sources using 1100 upward-going muons produced by neutrino interactions in the rock below and inside the MACRO detector in the underground Gran Sasso Laboratory. These data show no evidence for a possible neutrino point-like source or for possible correlations between gamma ray bursts and neutrinos. They have been used to set flux upper limits for candidate point-like sources which are in the range $10^{-14}-10^{-15}$ cm$^{-2}$ s$^{-1}$. ",
        "introduction": "\\label{sec:motiv}  The origin of cosmic rays is still largely an open question. The cosmic ray spectrum extends up to $10^{20}$ eV and the nature of the mechanisms capable of explaining such high energies is still unknown. Due to magnetic fields, charged cosmic particles are deflected from their original direction, hence the information on the position of their source is lost. On the other hand, protons of energies $\\gtrsim 10^{7}$ TeV and neutral particles, such as photons and neutrinos, point back to their sources since they are not deflected by magnetic fields. However, the universe should become opaque to protons with energies above $\\sim 5 \\cdot 10^{19}$ eV at distances of $\\sim 30$ Mpc due to photo-pion production when they interact with the Cosmic Microwave Radiation (CMBR) \\cite{GZK}. Photons are currently the main observation channel of our universe and the field of gamma-ray astronomy is now well established.  The idea of using neutrinos as probes of the deep universe was introduced in the sixties. Already in those years the first calculations on the diffuse neutrino flux from interactions of cosmic rays in the Galaxy \\cite{Greisen} and of high energy neutrinos from the Crab \\cite{Bahcall64} were performed.  Neutrinos are weakly interacting particles and are therefore much less absorbed than gamma rays, which are not only absorbed during their propagation, but can even be absorbed by the source producing them. Neutrinos can bring information on the deep interior of sources and on the far Universe. Several examples of detection of cosmic neutrinos already exist: solar neutrinos (from 0.1 MeV up to around 10 MeV), first detected by Homestake \\cite{Davis}, and neutrinos from SN1987A (from $\\sim 10$ MeV up to $\\sim 50$ MeV) detected by Kamiokande and IMB \\cite{Hirata87,Bionta87}. Nevertheless, neutrinos of astrophysical origin with energies larger than 100 MeV have not yet been observed. This observation would open the new field of high energy neutrino astronomy complementary to gamma ray astronomy. Moreover, an important hint on the existence of neutrino astronomy would come from the detection of photons of energies above 100 TeV. Such energies, in fact, cannot be explained by electron energy loss mechanisms (e.g. synchrotron radiation, bremsstrahlung and inverse Compton scattering) because electron acceleration is limited by the intense synchrotron radiation produced in the ambient magnetic fields. Therefore, alternate acceleration mechanisms which involve neutrinos are required.  Satellites, ground based imaging Cherenkov telescopes and extensive air shower arrays are currently investigating the universe, cosmic ray sources and acceleration mechanisms using photons as probes \\cite{Jackson93}. Current Space experiments typically work in the energy range up to about 30 GeV and ground-based terrestrial experiments have typical threshold energies of about 250 GeV. Ground based experiments have larger surfaces and longer time exposures than space experiments, therefore they can observe higher energies where fluxes are low. Nevertheless, they are limited at low energies by the large background due to gamma rays produced in the electromagnetic cascades induced by cosmic ray interactions in the atmosphere.  The EGRET detector on board the Compton Gamma Ray Observatory (CGRO) satellite has thus far furnished the largest amount of information on sources up to $\\sim 30$ GeV. The recent third EGRET catalogue \\cite{EGRET}, covering the observations made from 1991 to 1995, contains 271 sources observed with energies greater than 100 MeV. Between them, there are 5 pulsars, one probable radio galaxy (Cen A), 66 high-confidence identifications of blazars (BL Lac objects and radio quasars), 27 lower-confidence potential blazar identifications and a large number of identified supernova remnants (SNRs), and also 170 sources not yet identified firmly with known objects.  Satellite based detectors are providing observations on $\\gamma$-ray bursts (GRBs) capable of solving the mystery concerning their nature. The BATSE experiment \\cite{BATSE} on CGRO satellite has now detected more than 2500 $\\gamma$-ray bursts and the Italian-Dutch {\\it Beppo}SAX satellite \\cite{Frontera} is providing breakthroughs thanks to the precise measurement (the error box radius is at the level of 4') of the position of the bursts. Ground based experiments are looking for emissions above the TeV from GRBs: recently the Milagrito detector has found a correlation with the BATSE GRB970417a with chance probability $1.5 \\times 10^{-3}$ \\cite{Milagrito}.  Cherenkov telescopes at ground level such as the Whipple observatory, HEGRA, Cangaroo and University of Durham Mark 6 telescopes, have so far detected 8 sources emitting $\\gamma$-rays well above 300 GeV: the supernova remnants Crab \\cite{Lang}, Vela Pulsar (at a distance of only $\\sim 500$ pc)\\cite{Cangaroo0}, SN1006 \\cite{Cangaroo}, the extra-galactic BL Lac objects (highly variable active galactic nuclei) Mkn 421 ($z=0.031$)\\cite{Punch,Hegra}, Mkn 501 ($z < 0.034$) \\cite{Quinn}, PKS2155-304 ($z=0.116$) \\cite{Chadwick}, and the pulsars PSR1706-44 \\cite{Kifune}, PSR1259-63 \\cite{Sako}. The Whipple group detected the first source, the Crab supernova remnant \\cite{Lang}. The Crab is considered now as a standard candle for high energy gamma ray astronomy due to its gamma ray steady emission.  The number of sources detected so far by ground based experiments is much smaller than the number of sources detected by EGRET. One of the possible explanations is that high energy gamma rays are absorbed: TeV $\\gamma$-rays suffer absorption through pair production in intergalactic space on infrared light, PeV $\\gamma$s on the CMBR and EeV $\\gamma$s on radio-waves. This can even explain why the only BL Lac objects observed till now are also the nearest ones.  The discovery of TeV gamma rays emitted from the 8 sources quoted above shows the possibility of production from $\\pi^{o}$ decay and the possible existence of beam dump sources (see Sec.~\\ref{sec:astro}) producing high energy neutrinos. Nevertheless, ``a few TeV'' are energies not high enough to exclude a synchrotron radiation production mechanism. It could be completely excluded only for sources of energies above 100 TeV, but up to now, no source emitting at such energies has been discovered. There were some claims in the past, particularly about Cygnus X3, but they have not been confirmed \\cite{Hillas95}.  A different kind of neutrino production from astrophysical sources has been suggested by Gondolo and Silk \\cite{Gondolo}. If cold dark matter exists in the Galactic Center, it can be accreted by the black hole which probably is there. The cold dark matter is redistributed by the black hole into a cusp, which they call ``central spike''. If dark matter is made of neutral particles that can annihilate, such as the supersymmetric neutralino, the annihilation rate in the spike is strongly increased as it depends on the square of the matter density. Neutrinos can escape and produce relevant fluxes. For the neutralino, the fluxes are very high in the case of the presence of a central spike at the level of $10^{-15} \\div 10^{-14}$ cm$^{-2}$ s$^{-1}$ for $m_{\\chi} \\gtrsim 50$ GeV.  During the '70s and '80s the first generation underground detectors  of surface $\\sim 100$ m$^{2}$ have been measuring neutrinos. Those detectors were aimed at detecting proton decay for which atmospheric neutrinos were considered a background. Nevertheless, neutrinos were soon considered as an interesting signal themselves and results on searches for astrophysical sources of neutrinos were made. Previous results on the search for point-like sources have been published by the Kolar Gold Field experiment \\cite{Adarkar} and by the water Cherenkov detectors IMB \\cite{IMB} and Kamiokande \\cite{Kamiokande}. Other experiments (Baksan \\cite{Baksan} and AMANDA \\cite{AMANDA}) have presented preliminary results at conferences. MACRO has been detecting muon neutrinos since 1989 while it was still under construction. We present here the results of the search for astrophysical neutrino sources with MACRO during the period March 1989 - September 1999. ",
        "conclusions": "We have investigated the possibility that the sample of 1100 upward-going muons detected by MACRO since 1989 shows evidence of a possible neutrino astrophysics source. We do not find any significant signal with respect to the statistical fluctuations  of the background due to atmospheric neutrinos from any of the event directions or from any candidate sources. We also used the time information to look for correlations with gamma-ray bursts detected by BATSE and {\\it Beppo}SAX. Having found no excess of events with respect to the expected background we set muon and neutrino flux upper limits for point-like sources and for the cumulative search for catalogues of sources. These limits have been calculated taking into account the response of MACRO to various neutrino fluxes from candidate sources until energies $\\gtrsim 100$ TeV. These limits are for almost all of the considered sources the most stringent ones compared to other current experiments. They are about 1 order of magnitude higher than values quoted by most plausible neutrino source models except for the model in \\cite{Gondolo} which is seriously constrained.  {\\bf Acknowledgements}\\\\ We gratefully acknowledge the support of the director and of the staff of the Laboratori Nazionali del Gran Sasso and the invaluable assistance of the technical staff of the Institutions participating in the experiment. We thank the Istituto Nazionale di Fisica Nucleare (INFN), the U.S. Department of Energy and the U.S. National Science Foundation for their generous support of the MACRO experiment. We thank INFN, ICTP (Trieste) and NATO for providing fellowships and grants for non Italian citizens."
    },
    "0002/astro-ph0002171_arXiv.txt": {
        "abstract": "The calculations in Thomas-Fermi approximation show that in a gravitational field each cell of ultra dense matter inside celestial bodies obtains a very small positive electric charge. A celestial body is electrically neutral as a whole, because the negative electric charge exists at its surface. The positive volume charge is very small, on the order of magnitude it equals to $10^{-18}e$ per atom only. But it is sufficient to explain the occurrence of magnetic fields of the celestial bodies and the existence of a discrete spectrum of steady-state values of masses of planets, stars, and pulsars. ",
        "introduction": "According to the conventional point of view, gravity does not induce any electric polarization in the interior of celestial bodies and electric forces are never considered in the balance of matter of celestial bodies. Moreover, it is generally assumed that the electric interaction plays practically no role in astrophysics. It is a consequence of the comprehension that the appreciable electric polarization cannot arise in metals and other nonsegneto- and nonpiro-electric materials. It is entirely correct for all substances under action of small pressure. But, thus, one can disregard the fact that ultrahigh pressure transmutes all substances into plasma state and radically changes the properties of substance. In ultradense plasma, there is a different additional mechanism of the gravity-induced electric polarization.  In a large celestial body, consisting of ultradense plasma, this gravity-induced electric polarization (GIEP) can be rather great and can play a determining role in the formation of a number of features of the structure of a celestial body and its properties.  First of all, it concerns the following three problems, the statement and the solution of which change drastically:  - the distribution of pressure and density of matter inside a celestial body;  - the generation of a magnetic field by celestial bodies;  - the formation of a spectrum of steady-state values of masses of celestial bodies.  As a consequence, these features of the structure can influence the evolution of stars. ",
        "conclusions": "It is expedient  to underline the basic obtained results in summary.  1. The developed theory defines a concept of the steady-state values of masses of celestial bodies related to their equations of state and gives the possibility to calculate these values.  2. It gives the new way for the determination of the substance density distribution inside celestial bodies. According to early models, it was supposed that density of a substance inside celestial bodies grows more or less monotonically with depth and at the centre of a star, the density has the greatest value and even a black hole may exist there. According to the developed theory, the density of a substance inside a star is constant.  3. It is interesting to emphasize that the \"biography\" of such a star appears much poorer than in the Chandrasecar model.  There cannot exist a black hole inside a star, and it should not collapse with a temperature decrease. All the considered effects are based on an equilibrium of the Fermi system. Temperature does not influence the parameters of relativistic plasma. Therefore, a star with a mass close to the steady-state value (Eq.({\\ref{6050}})) is in a stable equilibrium not depending on temperature. The existing stars should retain the stability at any (even at zero) temperature. Therefore, a collapse of the already existing stars apparently is impossible. The instability of a star can arise with burning out of light nuclei - deuterium and helium - and with a related increasing of $\\beta$. This growth leads to the reduction of a steady-state value of mass (Eq.({\\ref{6050}})) and, probably, to the distraction of the stars wiht masses more than the steady-state value.  4. The developed theory determines the simple and essential mechanism of generation of the magnetic field by celestial bodies. All early models tried to solve the basic problem - to calculate the magnetic field of a celestial body. Such a statement of the basic problem of planetary and stellar magnetism is unacceptable at present. Space flights and a development of astronomy discovered a remarkable and earlier unknown fact: the magnetic moments of all celestial bodies are proportional to their angular momenta and the proportionality coefficient is determined by the ratio of world constants. The explanation of this phenomenon is the basic problem of planetary and stars magnetism nowadays. Early models cannot explain this phenomenon. The developed theory used for this explanation a standard mechanism.  5. It is possible to consider that now the predicted steady-state values of masses of celestial bodies and the predicted values of their magnetic moments are in a satisfactory agreement with the data of observations. But it is tempting to obtain these data to closer limit of accuracy. Two arrows in the upper part of Fig.5 mark masses of stars consisting of extremely light and heavy nuclei. These values are obtained from Eq.({\\ref{6050}}) without the use of any fitting parameters. In agreement wiht the developed theory, if stars have a \"usual\" chemical composition, there must be no stars outside of this interval (or these stars should be unstable). The histogram on Fig.5 is somewhat wider. It is interesting to understand, whether there is a principal deviation from the developed theory or it is a result of measuring errors. First of all, it requires a more careful and precise measurement of masses of binary stars."
    },
    "0002/astro-ph0002347_arXiv.txt": {
        "abstract": "Observations of synchrotron radiation across a wide range of wavelengths provide clear evidence that electrons are accelerated to relativistic energies in supernova remnants (SNRs). However, a viable mechanism for the pre--acceleration of such electrons to mildly relativistic energies has not yet been established. In this paper an electromagnetic particle--in--cell (PIC) code is used to simulate acceleration of electrons from background energies to tens of keV at perpendicular collisionless shocks associated with SNRs. Free energy for electron energization is provided by ions reflected from the shock front, with speeds greater than the upstream electron thermal speed. The PIC simulation results contain several new features, including: the acceleration, rather than heating, of electrons via the Buneman instability; the acceleration of electrons to speeds exceeding those of the shock--reflected ions producing the instability; and strong acceleration of electrons perpendicular to the magnetic field. Electron energization takes place through a variety of resonant and non--resonant processes, of which the strongest involves stochastic wave--particle interactions. In SNRs the diffusive shock process could then supply the final step required for the production of fully relativistic electrons. The mechanisms identified in this paper thus provide a possible solution to the electron pre--acceleration problem. ",
        "introduction": "The detection of radio synchrotron emission from shell--type supernova remnants (SNRs) is a clear indication that electrons of typically GeV energy are being accelerated in such objects. There is now convincing evidence that synchrotron emission from some remnants extends to X--ray wavelengths (Pohl \\& Esposito 1998): this implies the presence of electrons with energies of order 10$^{14}$eV. The prime example of this is the remnant of SN1006, recent observations of which using the ASCA (Koyama et al. 1995) and ROSAT (Willingale et al. 1996) spacecraft show that X--ray emission from the bright rim has a hard, approximately power--law spectrum. In contrast, emission from the centre is softer, with a strong atomic line component. The sharp edges and strong limb brightening observed at both X--ray and radio wavelengths indicate that: the acceleration site is the strong outer shock bounding the remnant; the acceleration is continuous; and the local diffusion coefficient of electrons near the shock front is substantially reduced relative to that in the general interstellar medium (Achterberg et al. 1994). The possibility that the X--ray emission from SN1006 is thermal bremsstrahlung has been examined by Laming (1998), and found to be less tenable than the synchrotron interpretation.  There is thus extensive observational evidence that the strong collisionless shocks bounding shell--type SNRs accelerate electrons to relativistic energies. The standard interpretation of extragalactic radio jet observations is also based on the premise that the relativistic electrons responsible for observed synchrotron emission are produced by shocks, although in this case the shock parameters are much less certain. Heliospheric shocks, on the other hand, do not generally appear to be associated with strong electron acceleration, perhaps because the Mach numbers of such shocks are much lower than those of SNRs and extragalactic radio jets, although Anderson et al. (1979) have published data showing that keV electrons are produced in the vicinity of the perpendicular bow shock of the Earth.  While diffusive shock acceleration (Axford et al. 1977; Krymsky 1977; Bell 1978; Blandford \\& Ostriker 1978) provides an efficient means of generating highly energetic electrons from an already mildly relativistic threshold, and can operate at oblique shocks as well as parallel ones (Kirk \\& Heavens 1989), the ``injection'' or ``pre--acceleration'' question remains very open: by what mechanisms can electrons be accelerated from background (sub--relativistic) energies to mildly relativistic energies (Levinson 1996)? In this paper, we investigate one of the possible answers, which has attractive ``bootstrap'' characteristics. Specifically, we suggest that waves are excited by collective instability of the non--Maxwellian population of ions reflected from a perpendicular shock front, and that these waves damp on thermal electrons, thereby accelerating them. Such a process was first proposed as a candidate acceleration mechanism for cosmic ray electrons by Galeev (1984). This model was developed further by Galeev et al. (1995), with the added ingredient of macroscopic electric fields implied by the need to maintain quasi--neutrality in a plasma with an escaping population of electrons. McClements et al. (1997) carried out a primarily analytical study of electron acceleration by ion--excited waves at quasi--perpendicular shocks, which was necessarily restricted to quantifying linear r\\'egimes of wave excitation and particle acceleration, in relation to inferred shock parameters.  Instabilities driven by shock--reflected ions at SNR shocks have also been invoked by Papadopoulos (1988) and Cargill \\& Papadopoulos (1988) as mechanisms for electron heating, rather than electron acceleration. On the basis of a simple analytical calculation, Papadopoulos predicted that strong electron heating would occur at quasi--perpendicular shocks with ``superhigh'' Mach numbers (specifically, shocks with fast magnetoacoustic Mach numbers $M_F > 30$--40) through the combined effects of Buneman (two--stream) and ion acoustic instabilities. In this model the Buneman instability, driven by the relative streaming of shock--reflected ions and upstream electrons, heats the electrons to a temperature $T_e$ much greater than the ion temperature $T_i$: in these circumstances, the ion acoustic instability can be driven unstable if there is a supersonic streaming between the electrons and either reflected or non--reflected (``background'') ions. Using a hybrid code, in which ions were treated as particles and electrons as a massless fluid, Cargill \\& Papadopoulos (1988) found that the electron heating predicted by Papadopoulos (1988) could occur in a self--consistently computed shock structure. However, as Cargill \\& Papadopoulos point out in the last paragraph of their 1988 paper, the use of a fluid model for the electrons means that hybrid codes cannot be used to investigate electron acceleration. Recently, Bessho \\& Ohsawa (1999) have used a particle--in--cell (PIC) code to investigate acceleration of electrons from tens of keV to highly relativistic energies at oblique shocks in which the electron gyrofrequency $\\Omega_e$ exceeds the electron plasma frequency $\\omega_{pe}$.  An improved theoretical understanding of electron acceleration at shocks is desirable not only for intrinsic interest, but also to enable observations of synchrotron and inverse Compton emission to be related quantitatively to shock parameters. However, almost all work on particle acceleration has concentrated on ions. There are several reasons for this. First, upstream momentum and energy fluxes are dominated by ions, and the shock structure problem therefore reduces essentially to that of isotropizing the ion distribution. Second, much of our understanding is based on the use of hybrid codes, in which electrons are represented as a fluid: such codes cannot provide information on electron acceleration. However, the very fact that electron dynamics does not appear to be important for shock structure allows us to separate the two problems: prescribing the ion parameters using the results of hybrid code simulations, we can examine in detail physical processes occurring on electron timescales. This is the approach followed in this paper. We describe the results of a fully nonlinear investigation, carried out by large scale numerical simulation using a PIC code and backed up by analytical and numerical studies, of the underlying plasma physics mechanisms. We consider the case $\\omega_{pe} > \\Omega_e$, which is qualitatively distinct from the strongly--magnetized r\\'egime investigated by Bessho \\& Ohsawa (1999). Our primary goal is finding a mechanism capable of producing mildly relativistic electrons: once they have attained rigidities comparable to those of shock--heated protons, they can undergo resonant scattering, and subsequent acceleration to relativistic energies can then proceed via the diffusive shock mechanism. Our approach enables us to test earlier predictions of both electron acceleration (Galeev 1984; Galeev et al. 1995; McClements et al. 1997) and electron heating (Papadopoulos 1988; Papadoulos \\& Cargill 1988) at very high Mach number astrophysical shocks. Simulation results are presented for a range of reflected ion speeds in Sect. 2; plasma instabilities occurring in the simulations, and other processes likely to play a role in electron acceleration and heating at SNR shocks, are identified in Sect. 3; and the results of these investigations are discussed in Sect. 4. ",
        "conclusions": ""
    },
    "0002/astro-ph0002037_arXiv.txt": {
        "abstract": "We consider gravitational clustering from primoridal non-Gaussian fluctuations provided by a $\\chi^2$ model, as motivated by some models of inflation.  The emphasis is in signatures that can be used to constrain this type of models from large-scale structure galaxy surveys. Non-Gaussian initial conditions provide additional non-linear couplings otherwise forbidden by symmetry that cause non-linear gravitational corrections to become important at larger scales than in the Gaussian case. In fact, the lack of hierarchical scaling in the initial conditions is partially restored by gravitational evolution at scales $k> 0.1$ h/Mpc.  However, the bispectrum shows much larger amplitude and residual scale dependence not present in evolution from Gaussian initial conditions that can be used to test this model against observations.  We include the effects of biasing and redshift distortions essential to compare this model with galaxy redshift surveys. We also discuss the effects of primordial non-Gaussianity on the redshift-space power spectrum and show that it changes the shape of the quadrupole to monopole ratio through non-linear corrections to infall velocities. ",
        "introduction": "The current paradigm for the formation of large-scale structures in the universe is that small primordial fluctuations, with a roughly scale-invariant spectrum, are amplified by non-linear gravitational instability to form the galaxies, clusters and superclsuters that we see today in galaxy surveys.  A common assumption is that the statistics of these primordial fluctuations is Gaussian, which is consistent with current observations.  However, at this point we cannot rule out the possibility that primordial fluctuations were in fact non-Gaussian to some extent.  Fortunately, rapid progress in microwave background anisotropy experiments and large galaxy redshift surveys in the next few years will provide high quality data that can be used to test the nature of primordial fluctuations.   The consideration of non-Gaussian initial conditions is a difficult issue because there is an infinite class of non-Gaussian models.  In a Gaussian field $\\phi$, $n-$point connected correlation functions, $\\xi_n = \\lexp \\phi_{1} \\ldots \\phi_{n} \\rexp_{c}$ of order $n \\geq 3$ vanish; however, in a non-Gaussian field higher-order correlations can virtually behave arbitrarily subject to the constraint of realizability, that is, that the multivariate probability distributions be positive definite.  In general, however, we can divide non-Gaussian models into {\\em weakly} and {\\em strongly} non-Gaussian, depending on the magnitude of normalized cumulants $s_{p} \\equiv \\lexp \\d^p \\rexpc / \\lexp \\d^2 \\rexp^{p/2}$ of the smoothed density contrast $\\d$ compared to unity.  A minimal way of achieving strongly non-Gaussian models is by {\\em dimensional scaling}, where the hierarchy of $n$-point correlation functions obeys $\\xi_n \\propto \\xi_2^{n/2}$.  In this case, normalized cumulants are numbers not necessarily smaller than unity.  This is opossed to {\\em weakly non-Gaussian models} where normalized cumulants are forced to be small quantities, e.g.  {\\em hierarchical scaling} models, where $\\xi_n \\propto \\xi_2^{n-1}$ with $\\xi_2 \\ll 1$.  In this paper we consider $\\chi^2$ initial conditions, which belong to the class of dimensional scaling models. As a strongly non-Gaussian model, it has the potential to be more easily constrained or detected by large-scale structure observations than primordial weakly non-Gaussian models. Furthermore, there is a number of inflationary models in the literature that motivate $\\chi^2$ initial conditions (e.g. Kofman et al. 1989; Antoniadis et al. 1997, Linde \\& Muhanov 1997; Peebles 1997).  In addition, this model has been recently argued to fit a significant set of observational constraints (Peebles 1999a, 1999b). It is also possible that this particular model may be a good representation of the general behavior of dimensional scaling models, and thus provide general insight about their advantages and disadvantages. Furthermore, models in which primordial fluctuations are generated by topological defects also generally obey dimensional scaling (e.g. Turok \\& Spergel 1991).   Previous work on clustering from non-Gaussian initial conditions was done using numerical simulations (e.g. Moscardini et al. 1991; Weinberg and Cole 1992; Coles et al. 1993; White 1999), and perturbation theory (Fry \\& Scherrer 1994; Jaffe 1994; Chodorowski \\& Bouchet 1996; Gazta\\~naga \\& Fosalba 1998). In this work, we concentrate on aspects of $\\chi^2$ initial conditions that can be tested with current and future observations in galaxy redshift surveys. In particular, we concentrate on the study of the redshift-space power spectrum and bispectrum; applications of these results to observations will be considered elsewhere (Scoccimarro et al. 2000).   Different aspects of the effects of primordial non-Gaussianity on the power spectrum have been considered in the literature (e.g. Feldman, Kaiser \\& Peacock 1994; Stirling \\& Peacock 1996; Sutherland et al. 1999). The skewness of the smoothed density field in texture models was studied by Gazta\\~naga \\& M\\\"ah\\\"onen (1996), whereas the impact of hierarchical scaling models of primordial non-Gaussianity on the bispectrum was considered by Verde et al. (1999). Our approach is complementary to recent studies of the impact of primordial non-Gaussian models in other aspects of large-scale structure (e.g. Robinson, Gawiser \\& Silk 1999; Koyama, Soda \\& Taruya 1999; Willick 1999; Pierpaoli, Garc\\'{\\i}a-Bellido \\& Borgani 1999).   This paper is organized as follows. In Section~2 we review perturbation theory results regarding the evolution of the bispectrum from general initial conditions. Section~3 presents results for the particular case of $\\chi^2$ initial conditions. In Section~4 we discuss the effects of galaxy biasing and Section~5 those of redshift distortions. Finally, we present our conclusions in Section~6. ",
        "conclusions": "We have explored the predictions of $\\chi^2$ initial conditions for clustering statistics, in particular the power spectrum and bispectrum. We found that extrapolation of the initial conditions using linear perturbation theory is not accurate enough at the scales best probed by current and future galaxy redshift surveys ($k \\ga 0.05$ h/Mpc). This is not surprising, non-Gaussian initial conditions provide additional non-linear couplings otherwise forbidden by symmetry, thus non-linear gravitational corrections can become important at larger scales than in the Gaussian case.   We showed that when non-linear corrections are included, the bispectrum shows a similar configuration dependence as in the Gaussian case; however its amplitude is much larger than the latter and has a residual scale dependence not present in evolution from Gaussian initial conditions. We included the effects of galaxy biasing and showed that the galaxy bispectrum obeys the same relation to the dark matter bispectrum as in the Gaussian case with a proper redefinition of the non-linear bias parameter. Thus, a large linear bias decreases the dependence of the bispectrum on triangle configuration, and a non-linear bias contributes a constant independent of configuration just as in the Gaussian case. The effects of redshift distortions were shown to change the overall amplitude of Gaussian and non-Gaussian bispectra, but clear difference between them remains.  Thus, we conclude that the bispectrum is a very useful statistic to probe the Gaussianity of initial conditions, at least for models where the scaling is dimensional as in the $\\chi^2$ case. Application of these results to current galaxy surveys will be reported elsewhere (Scoccimarro et al. 2000).   We also discussed the effects of non-Gaussianity on the power spectrum in redshift space. The shape of the quadrupole to monopole ratio is very sensitive to the large-scale pairwise velocity which depends on the primordial bispectrum through non-linear corrections. For the $\\chi^2$ model, this leads to a smaller infall velocity and thus a supression of overall power and smaller quadrupole zero crossing scale compared to the case of Gaussian initial conditions. This provides an additional signature of primordial non-Gaussianity that can be explored with the future generation of galaxy redshift surveys."
    },
    "0002/astro-ph0002201_arXiv.txt": {
        "abstract": "We report the detection with \\chandra ~of a Low-Luminosity AGN (LLAGN) in the Low Ionization Emission Line Region (LINER) hosted by Hydra~A, a nearby ($z$=0.0537) powerful FR~I radio galaxy with complex radio and optical morphology. In a 20 ks ACIS-S exposure during the calibration phase of the instrument, a point source is detected at energies $\\gtrsim$ 2 keV at the position of the compact radio core, embedded in diffuse thermal X-ray emission ($kT \\sim 1$ keV) at softer energies. The spectrum of the point source is well fitted by a heavily absorbed power law with intrinsic column density N$_H^{int} \\sim 3 \\times 10^{22}$ \\nh~ and photon index $\\Gamma \\sim 1.7$. The intrinsic (absorption-corrected) luminosity is $L_{2-10~keV} \\sim 1.3 \\times 10^{42}$ \\lum. These results provide strong evidence that an obscured AGN is present in the nuclear region of Hydra~A. We infer that the optical/UV emission of the AGN is mostly hidden by the heavy intrinsic reddening. In order to balance the photon budget of the nebula, we must either postulate that the ionizing spectrum includes a UV bump or invoke and additional power source (shocks in the cooling flow or interaction with the radio jets).  Using an indirect estimate of the black hole mass and the X-ray luminosity, we infer that the accretion rate is low, suggesting that the accretion flow is advection dominated. % Finally, our results support current unification schemes for radio-loud sources, in particular the presence of the putative molecular torus in FR~Is.  These observations underscore the power of the X-rays and of \\chandra~in the quest for black holes. ",
        "introduction": "Recent {\\it HST} and ground-based optical observations provided strong evidence that many nearby galaxies harbor supermassive black holes (e.g., Kormendy \\& Richstone 1995), possibly accreting at sub-Eddington luminosities (Fabian \\& Rees 1995). About 40\\% of nearby early-type galaxies exhibit signs of mild nuclear activity, in the form of weak non-thermal radio cores (Sadler et al. 1989) and Low Ionization Emission Lines (LINERs; Heckman 1986; Ho, Fillippenko, \\& Sargent 1997a). These results collectively suggest that many nearby galaxies may harbor weak nuclear activity in the form of a Low-Luminosity Active Galactic Nucleus (LLAGN; e.g., Ho 1999a).  Thanks to their high penetrating power, X-rays provide an optimal window to search for weak nuclear activity.  Indeed, previous X-ray \\rosat~ images of a handful of galaxies show the presence of a central unresolved nucleus within the 5\\arcsec~ HRI resolution (e.g., Fabbiano 1996). Indirect clues are provided by \\asca~ spectral constraints: a heavily absorbed power law component is often measured at energies $\\gtrsim 2$ keV, with intrinsic luminosities L$_{2-10~keV} \\sim 10^{40-42}$ \\lum~, photon indices $\\Gamma_{2-10~keV} \\sim 1.5-1.7$, and a narrow Fe emission line at 6--7 keV in a few cases, suggestive of a LLAGN (Makishima et al. 1994; Ptak et al. 1999; Sambruna, Eracleous, \\& Mushotzky 1999; Terashima et al. 1999). Within the coarse angular resolution of these detectors, alternative scenarios can not be ruled out in many cases (e.g., a starburst or X-ray binaries). Unambiguous evidence for nuclear activity would be provided by the detection of a point source at the galaxy center.  With its unprecedented angular resolution (0.5\\arcsec), wide-band coverage (0.2--10 keV), and high sensitivity, \\chandra~ is uniquely suited to this task.  In this Letter, using recent \\chandra~calibration observations we discover a LLAGN in the LINER harbored by the nearby cD galaxy Hydra~A ($z$=0.0537).  Host of the powerful FR~I radio source 3C~218, Hydra~A is the dominant member of the poor cluster of galaxies A780, and famous for its twin-jet radio morphology (Taylor et al. 1990). The nuclear region contains a $\\sim$ 6\\arcsec~emission-line nebula (Baum et al. 1989; Heckman et al. 1989) and a disk of star formation (McNamara 1995; Melnick, Gopal-Krishna, \\& Terlevich 1997). At the position of the nucleus, a LINER-like optical and UV spectrum is observed (Hansen, J{\\o}rgensen, \\& N{\\o}rgaard-Nielsen 1995), which, together with the powerful lobe radio emission, led to the classification of Hydra~A as a Weak Line Radio Galaxy (Tadhunter et al. 1998). The X-ray emission from Hydra~A in previous {\\it Einstein}, \\rosat, and \\asca~ images is dominated by the cluster and the inner ($\\lesssim$ 1.5\\arcmin) cooling flow, with $L^{cluster}_{0.5-4.5~keV} \\sim 2 \\times 10^{44}$ \\lum~ (David et al. 1990; Ikebe et al. 1997).  In the following, we assume a Friedman cosmology with $H_0=75$ km s$^{-1}$ Mpc$^{-1}$ and $q_0=0.5$. At the luminosity distance of Hydra~A (217.46 Mpc), 1\\arcsec=0.95 kpc. ",
        "conclusions": "Hydra~A is the second powerful radio galaxy observed with \\chandra, and the second instance where a nuclear X-ray point source is detected. A luminous AGN was also discovered in an ACIS-S 20 ks image of the FR~II radio galaxy 3C~295 (Harris et al. 2000), with $\\Gamma \\sim 1.8$ and $L_{0.2-10~keV} \\sim 7 \\times 10^{43}$ \\lum. As in Hydra~A, the AGN in 3C~295 resides in a cooling flow. These results confirm that optically-weak, narrow-emission line radio sources harbor true AGNs as their brighter, broad-emission line Seyfert-like counterparts. Thanks to \\chandra, a detailed study of the central engines of these systems becomes possible for the first time.  The large X-ray column density we measure in the nucleus of Hydra~A, N$_H^{int} \\sim 3 \\times 10^{22}$ \\nh, implies an optical extinction $A_V=10$, assuming galactic gas-to-dust ratios (Bohlin, Savage, \\& Drake 1978). This means that the optical and UV continuum from the AGN are strongly (factor $\\gtrsim 300$) suppressed, accounting for why the AGN was not previously detected at these wavelengths. The weak UV continuum detected in archival IUE spectra (Hansen et al. 1995 and our own inspection of an unpublished spectrum) cannot be due to the AGN; most likely candidates are the starburst and/or the cooling flow. Future {\\it HST} observations should detect only an extended thermal component to the UV light.  We now turn to the issue of whether the LLAGN is capable of photoionizing the emission-line nebula identified with the nucleus (Baum et al. 1989; Heckman et al.  1989). To make this assessment we first evaluate the total reddening to the nuclear emission-line source, $E_{\\rm B-V}$. Since the nuclear spectrum resembles that of LINERs, we assume that the intrinsic Balmer decrement is H$\\alpha$/H$\\beta=3.1$ (e.g., Ho, Filippenko, \\& Sargent 1997b), leading to an estimate of $E_{\\rm B-V}\\approx 0.33$\\footnote{This is considerably higher than $E_{\\rm B-V}=0.15$ given by Hansen et al. (1995) based on the Ly$\\alpha$/H$\\beta$ ratio. Hansen et al. have probably underestimated the reddening because the Ly$\\alpha$ flux was measured in the large IUE aperture which very likely includes a contribution from the circumnuclear starburst region.}.  With this value of the reddening and the H$\\beta$ flux measured by Hansen et al. (1995) we obtain the rate of emission of H$\\beta$ photons from the nebula as $Q_{\\rm H\\beta}=4.1\\times 10^{51}$~s$^{-1}$. If the H$\\beta$ emission is the result of case~B recombination, then atomic physics implies $Q_{\\rm H\\beta}=0.12\\; Q_{\\rm ion}$ (Osterbrock 1989), where $Q_{\\rm ion}$ is the ionizing photon rate from the active nucleus. Thus, the observed H$\\beta$ flux requires that $Q_{\\rm ion}=(3.4\\times 10^{52}/f_{\\rm c})$~s$^{-1}$, where $f_{\\rm c}$ is the covering fraction of the source by the nebula. In contrast, the observed ionizing photon rate, assuming that the X-ray power-law spectrum extends all the way through the UV band, is $Q_{\\rm ion}=1.1\\times 10^{52}$~s$^{-1}$, which falls short of the required rate by a considerable factor, even if $f_{\\rm c}=1$! In fact, if we model the spectral energy distribution (SED) of Hydra A after those observed in LINERs and other LLAGN, some of which are radio-loud (Ho 1999b), the problem persists.  In order to balance the photon budget of the nebula, we must either postulate that the ionizing spectrum includes a UV bump or invoke and additional power source. If we assume there is a UV bump, we can parameterize it following Mathews \\& Ferland (1987) and normalize the SED according to the \\chandra~ X-ray spectrum. We then find an ionizing photon output rate of $Q_{\\rm ion}=6.4\\times 10^{53}$~s$^{-1}$, which is more than enough to power the emission lines. On the other hand, one or more other power sources may contribute to the ionization of the nebula.  X-rays from shocks in the cooling flow are plausible (e.g., Heckman et al. 1989). Another possibility is mechanical interaction of the emission-line gas with the radio jets, which is particularly relevant since the images of Hansen et al. (1995) show superpositions of radio and emission-line knots.  To get a handle on the nature of the accretion flow, we compared the X-ray luminosity to the limiting Eddington luminosity through a rough estimate of the central black hole mass in Hydra~A, $M_{BH}$.  The apparent $V$ magnitude of the host galaxy in Hydra~A, $m_V=13.7$ (Sandage 1973), and Figure 8 of Magorrian et al. (1998) imply $M_{BH} \\sim 4 \\times 10^9 M_{\\odot}$, consistent with the values dynamically measured in other giant ellipticals (e.g., Ford et al. 1994).  The Eddington luminosity is thus $L_{Edd} \\sim 5 \\times 10^{47}$ \\lum. Assuming $L_X=0.1 L_{total}$, the luminosity relative to Eddington is $L_{total}/L_{Edd} \\sim 2 \\times 10^{-5}$. This places Hydra~A in the regime of an Advection-Dominated Accretion Flow (ADAF; e.g., Narayan, Mahadevan, \\& Quataert 1998), similar to other WLRGs (Sambruna et al. 1999). A powerful diagnostic of the accretion flow structure in Hydra~A will be afforded by future higher-sensitivity X-ray observations, such as delivered by {\\it XMM}. The EPIC spectrum will allow us to study in more detail the nuclear X-ray properties, in particular whether an Fe line is present at 6--7 keV. The line energy and profile could allow us to discriminate between a standard Seyfert-like disk and an ADAF (e.g., Sambruna \\& Eracleous 1999).   Finally, we note that the nuclear X-ray absorption in Hydra~A is much larger (factor 10) than the optical/UV extinction to the emission-line nebula (see above), and consistent with the column measured in the radio toward the core (Taylor 1996). This suggests that the X-ray absorber lies further in than the emission-line regions, close to the central black hole. A possible candidate is the molecular torus on parsec scales postulated by unification models (Urry \\& Padovani 1995). Indeed, the radio and optical observations indicate an edge-on geometry for Hydra~A (Taylor et al. 1990; Baum et al. 1989), such that the torus intercepts the line of sight to the nucleus.  Our results thus support current unification models for radio-loud sources and in particular the presence of an obscuring torus in FR~I radio galaxies. Clearly, large unbiased samples are needed to reach firmer conclusions, which we anticipate from \\chandra~in the next few years."
    },
    "0002/astro-ph0002084_arXiv.txt": {
        "abstract": "We report on the X-ray spectral and timing behavior of the recurrent X-ray transient 4U~1630--47 for 51 \\it RXTE \\rm observations made during the decay of its 1998 outburst.  The observations began when the source was still relatively bright, and, during one of the early observations, a QPO with a non-Lorentzian profile occurred near 6~Hz. As the source decayed, the X-ray flux dropped exponentially with an e-folding time of 14.4~d.  The exponential decay was interrupted by an increase in the X-ray flux, and a secondary maximum occurred 89~d after the onset of the outburst.  A transition marked by significant changes in the timing and spectral properties of the source occurred 104~d after the start of the outburst.  The transition is similar to soft-to-hard state transitions observed in other black hole candidate X-ray binaries.  Most of the changes associated with the transition occurred in less than 2~d.  The timing changes include an increase in the continuum noise level from less than 4\\% RMS to greater than 10\\% RMS and the appearance of a quasi-periodic oscillation (QPO) at 3.4~Hz with an RMS amplitude of 7.3\\% in the 2-21~keV energy band.  At the transition, the energy spectrum also changed with an abrupt drop in the soft component flux in the \\it RXTE \\rm band pass.  A change in the power-law photon index from 2.3 to 1.8, also associated with the transition, occurred over a time period of 8~d.  After the transition, the source flux continued to decrease, and the QPO frequency decayed gradually from 3.4~Hz to about 0.2~Hz. ",
        "introduction": "Although the recurrent X-ray transient and black hole candidate (BHC) 4U~1630--47 has been studied extensively since its first detected outburst in 1969 (\\cite{priedhorsky86}), interest in this source has intensified due to observations made during its 1998 outburst.  During the 1998 outburst, radio emission was detected for the first time (\\cite{hjellming99}).  Although the source was not resolved in the radio, the optically thin radio emission suggests the presence of a radio jet.  Also, low frequency quasi-periodic oscillations (QPOs) were discovered during the 1998 outburst (\\cite{dieters98a}) using the \\it Rossi X-ray Timing Explorer \\rm (\\it RXTE\\rm).  Here, we report on X-ray observations of 4U~1630--47 made with \\it RXTE \\rm (\\cite{brs93}) during the decay of its 1998 outburst.  We compare the X-ray light curve to those of other BHC X-ray transients and study the evolution of the spectral and timing properties during the decay.  Like many other X-ray transients, the light curve of 4U~1630--47 shows an exponential decay and a secondary maximum (\\cite{csl97}). During the early part of the decay, when the X-ray flux was high, 4U~1630--47 showed canonical soft state characteristics (\\cite{v95};~\\cite{nowak95};~\\cite{ct96}), including an energy spectrum with a strong soft component and a steep power-law and relatively low timing variability with a fractional RMS (Root-Mean-Square) amplitude of a few percent.  Later in the decay, we observe a transition to a spectrally harder and more variable state, which has similarities to transitions observed for GS~1124--68 (\\cite{ebisawa94};~\\cite{miyamoto94}) and GRO~J1655--40 (\\cite{mendez98}) near the ends of their outbursts.  In this paper, we describe the 4U~1630--47 X-ray light curve for the 1998 outburst and the \\it RXTE \\rm observations (\\S 2).  In \\S 3 and \\S 4, we present results of modeling the power and energy spectra, respectively.  In \\S 5, we examine the transition in more detail, and \\S 6 contains a discussion of the results.  Finally, \\S 7 contains a summary of our findings. ",
        "conclusions": "We have analyzed data from 51 \\it RXTE \\rm observations of 4U~1630--47 during the decay of its 1998 outburst to study the evolution of its spectral and timing properties.  During the decay, the X-ray flux dropped exponentially with an e-folding time of about 14.4~d, which is short compared to most other BHC X-ray transients.  The e-folding time was nearly the same (14.9~d) for the decay of the 1996 outburst, which may indicate that this time scale is set by some property of the system that does not change between outbursts.  For the 1998 outburst, the decay was interrupted by a secondary maximum, which is commonly observed for BHC X-ray transients.  Our analysis of the 4U~1630--47 power spectra indicates that 0.2~Hz to 11~Hz QPOs with RMS amplitudes between 2\\% and 9\\% occurred during the observations. During one of our early observations, when the source was relatively bright, a QPO occurred near 6~Hz with a profile that cannot be described by a single Lorentzian.  Similar asymmetric QPO peaks have been observed previously for GS~1124--68 (\\cite{belloni97}) and XTE~J1748--288 (Revnivtsev et al.~1999). For all three sources (4U~1630--47, GS~1124--68 and XTE~J1748--288), the QPO is well-described by a combination of two Lorentzians.  Near the end of the outburst, an abrupt change in the 4U~1630--47 spectral and timing properties occurred, and we interpret this change as evidence for a soft-to-hard state transition.  Our observations indicate that most of the changes in the emission properties, associated with the transition, occurred over a time period less than 2~d.  The timing properties changed after the transition with an increase in the continuum noise level and the appearance of a QPO.  A 3.4~Hz QPO appeared immediately after the transition, and, in subsequent observations, the QPO frequency decreased gradually to about 0.2~Hz.  At the transition, the energy spectrum also changed with an abrupt drop in the soft component flux in the \\it RXTE \\rm band pass, which was probably due to a drop in the inner disk temperature.  A change in the power-law photon index from 2.3 to 1.8, also associated with the transition, occurred over a time period of 8~d.  Although many of these changes are typical of soft-to-hard state transitions, the QPO behavior and the short time scale for the transition are not part of the canonical picture for state transitions (\\cite{v95};~\\cite{nowak95};~\\cite{ct96}). Finally, we note that 4U~1630--47 exhibits interesting behavior (e.g., state changes and QPOs) below a flux level of $10^{-9}$~erg~cm$^{-2}$~s$^{-1}$, indicating that observing programs for X-ray transients should be designed to follow these sources to low flux levels."
    },
    "0002/astro-ph0002421_arXiv.txt": {
        "abstract": "\\small We discuss the composition of dust and ice along the line of  sight to the Galactic Center (GC) based on analysis of  mid-infrared spectra (2.4--13 $\\mu$m) from the Short  Wavelength Spectrometer on the Infrared Space Observatory  (ISO).  We have analyzed dust absorption features arising in the  molecular cloud material and the diffuse interstellar medium  along the lines of sight toward Sagittarius A* and the  Quintuplet sources, GCS3 and GCS4. It is evident from the  depth of the 3.0 $\\mu$m H$_2$O and the 4.27 $\\mu$m CO$_2$  ice features that there is more molecular cloud material  along the line of sight toward Sgr A* than GCS3 and 4. In  fact, Sgr A* has a rich infrared ice spectrum with evidence  for the presence of solid \\methane, NH$_3$, and  possibly HCOOH.  Hydrocarbon dust in the diffuse interstellar medium along  the line of sight to the GC is characterized by absorption  features centered at 3.4 $\\mu$m, 6.85 $\\mu$m, and 7.25  $\\mu$m. Ground-based studies have identified the 3.4 $\\mu$m  feature with aliphatic hydrocarbons, and ISO has given us the first meaningful observations of the corresponding modes  at longer wavelengths. The integrated strengths of these  three features suggest that hydrogenated  amorphous carbon is their carrier. We attribute an absorption feature centered at 3.28 $\\mu$m in the GCS3  spectrum to the C-H stretch in aromatic hydrocarbons.  This feature is not detected, and its C-C stretch counterpart appears to be weaker, in the \\sgra\\ spectrum. One  of the key questions which now arises is whether   aromatics are a widespread component of the diffuse  interstellar medium, analogous to aliphatic  hydrocarbons. ",
        "introduction": "Bright infrared (IR) sources combined with 30 magnitudes of visual extinction have made the Galactic Center (GC) an attractive line of sight for the  study of interstellar dust composition. Absorption features in the spectra of GC IR sources have been used extensively to characterize interstellar grain material, including  silicates \\nocite{becklin_etal78I,roche_aitken85} (Becklin \\etal\\ 1978; Roche \\& Aitken 1985) and  hydrocarbons \\nocite{willner_etal79,mcfadzean_etal89,sandford_etal91,pendleton_etal94} (Willner \\etal\\ 1979; McFadzean \\etal\\ 1989; Sandford \\etal\\ 1991; Pendleton \\etal\\ 1994). Previous observations of the  GC suggested the presence of absorptions expected in icy grain material, most notably the 3$\\mu$m band of \\water-ice (e.g., McFadzean \\etal\\ 1989), but until recently, the explanation that one or more molecular clouds obscure the GC and contribute to its total extinction had not been completely accepted since the 3$\\mu$m feature has a substantially different profile than that observed in dense cloud sources (Tielens \\etal\\ 1996; Wada \\etal\\ 1991). The presence of molecular cloud material in the GC was confirmed via detection  of solid CO$_2$ feature by the Short Wavelength Spectrometer (SWS) on the Infrared Space Observatory (ISO) \\nocite{gerakines_etal99,degraauw_etal96b,lutz_etal96} (Gerakines \\etal\\ 1999; de Graauw \\etal\\ 1996b; Lutz \\etal\\ 1996).  Solid organic material is a well-established component of the diffuse interstellar medium.    The complex 3.4 \\micron\\ feature, ubiquitous along sightlines which sample the diffuse ISM in our Galaxy and other galaxies \\nocite{pendleton_etal94,bridger_wright_geballe93}(e.g. Pendleton \\etal\\ 1994, Bridger, Wright, \\& Geballe 1993),  is identified with short-chained aliphatic hydrocarbons with a methylene/methyl (-CH$_2$-/-CH$_3$) group ratio between 2--2.5 (Sandford \\etal\\ 1991; Pendleton \\etal\\ 1994). The CH deformation mode at  6.8 \\micron\\ was detected in KAO data (Tielens \\etal\\ 1996), although the weaker 7.3 \\micron\\ band could not be confirmed at the time.  One model proposes that organic refractory matter in the diffuse ISM is the product of energetic processing of simple ices (such as CO, \\water, \\methanol) in dense clouds \\nocite{allamandola_sandford_valero88,greenberg_etal95, strazzulla_castornia_palumbo95,moore_ferrante_nuth96} (e.g., Allamandola, Sandford, \\& Valero 1988;  Greenberg \\etal\\ 1995; Strazzulla, Castornia, \\& Palumbo 1995; Moore, Ferrante, \\& Nuth 1996). However, the discovery of spectroscopic similarity between diffuse interstellar and circumstellar organics has prompted a reassessment of the origin of organic  matter in the ISM \\nocite{lequeux_muizon90,chiar_etal98crl}(Lequeux \\& Jourdain de Muizon 1990; Chiar \\etal\\ 1998b).  The striking similarity of the circumstellar and interstellar 3.4 \\micron\\ features indicates that at least  some  of the organic carriers which lead to the 3.4 \\micron\\ absorption in the  diffuse ISM originate as stardust (Chiar \\etal\\ 1998b). The case for production of the 3.4 \\micron\\ band carrier in dense clouds and thus its presence in organic  refractory mantles is further weakened by the finding of \\nocite{adamson_etal99} Adamson \\etal\\ (1999) that, unlike the 9.7 \\micron\\ silicate feature \\nocite{aitken_etal86}(e.g., Aitken \\etal\\ 1986),  the 3.4 \\micron\\ feature along the line of sight toward \\sgra\\ (IRS7) is not polarized.   Hence,  the  3.4 \\micron\\ band is not produced by a carrier residing in a mantle on a silicate core, but rather by very small (unaligned) grains.  Nevertheless, energetic processing of ices in dense clouds may still play some role in the production of interstellar organics since diffuse ISM dust cycles into dense clouds fairly rapidly \\nocite{mckee89,jones_etal94}($\\sim 3\\times10^7$ years; McKee 1989; Jones \\etal\\ 1994).  Hydrogenated Amorphous Carbon (HAC) is  a possible candidate for the organic component of diffuse interstellar dust (Tielens \\etal\\ 1996; Pendleton \\etal\\ 1994), and several laboratory groups are attempting to produce HAC analogs with a -CH$_2$-/-CH$_3$ ratio which matches the observed feature \\nocite{furton_etal99,duley_etal98}(e.g., Furton \\etal\\ 1999; Duley \\etal\\ 1998).  Previous comparisons of HACs with the interstellar 3.4 \\micron\\ feature have shown that many  HACs produced in the laboratory are lacking in -CH$_3$ groups relative to -CH$_2$- groups, providing an unsatisfactory fit to the observed feature (e.g., Pendleton \\etal\\ 1994).  Recently, Duley \\etal\\ (1998) have shown that the 3.38 \\micron\\ -CH$_3$ subfeature is strengthened at lower temperatures (77\\,K), and the resulting HAC with -CH$_2$-/-CH$_3<1$, provides a better match to the 3.4 \\micron\\ interstellar feature, compared with higher temperature (e.g., 300\\,K) HACs.  Recently, a feature near 6.2 \\micron\\ was discovered along sightlines toward the GC Quintuplet sources and several dusty late-type WC Wolf-Rayet stars, and has been attributed to the aromatic C-C stretch in polycyclic aromatic hydrocarbon molecules (PAH) in the diffuse ISM along the line of sight (Schutte \\etal\\ 1998).   PAHs are a certain component of the carbonaceous material in the Galaxy, as shown by the dominance and ubiquitous presence of the  family of PAH emission features at 3.3, 6.2, 7.7, 8.6, and 11.3 \\micron\\  in diffuse emission in the Galactic disk \\nocite{giard_etal88,mattila_etal96,ristorcelli_etal94}(Giard \\etal\\ 1988;  Ristorcelli \\etal\\ 1994; Mattila \\etal\\ 1996).   Therefore, the detection of the aromatic C-C stretch at 6.2 \\micron\\ in diffuse ISM sightlines is not surprising; its carrier may be in the form of  gaseous polycyclic aromatic hydrocarbons (PAHs) or PAH clusters with sizes up to a few thousand C atoms \\nocite{schutte_etal98}(Schutte \\etal\\ 1998).  The corresponding C-H stretch at $\\sim3.3$ \\micron\\ was not detected in absorption in the diffuse ISM, although high signal-to-noise data, necessary to detect this weak feature, were not available for many sightlines.  However, an absorption feature at 3.25 $\\micron$ is detected in ground-based data in five heavily extincted molecular cloud sightlines \\nocite{brooke_sellgren_geballe99,sellgren_etal95}(Sellgren \\etal\\ 1995; Brooke, Sellgren, \\& Geballe 1999).  In this paper we discuss the composition of the  hydrocarbon dust component as well as the molecular cloud ices, along the line of sight to the Galactic Center based on 2.5--13 \\micron\\ spectra from the Infrared Space Observatory's Short Wavelength Spectrometer (ISO-SWS).  ISO-SWS had revealed the previously elusive 2.95 \\micron\\ \\ammonia\\ feature in dense cloud material along the GC line-of-sight.  These observations have also provided us with the first meaningful observations of hydrocarbon modes at long wavelengths (5--8\\micron) corresponding to the well-studied 3.4 \\micron\\ absorption feature, associated with the diffuse ISM, allowing us to narrow down the list of possible hydrocarbon candidates (\\S\\ref{aliphatics}).  We also present the first detection of aromatic absorption, at 3.28 \\micron, unassociated with molecular cloud dust. This paper is organized as follows:  In \\S\\ref{sources} we briefly discuss what is known about the nature of the sources, and in \\S\\ref{obs} we explain the data reduction, and present the ISO-SWS spectra of 3 sources in the Galactic Center.  In the following sections, we discuss the nature of the molecular cloud IR absorption features (\\S\\ref{molecular}) and hydrocarbon absorption features (\\S\\ref{diffuse}) along the line of sight.   Finally, we  discuss possible scenarios for the origin of hydrocarbons in the GC region (\\S\\ref{discussion}) and summarize our results (\\S\\ref{summary}). ",
        "conclusions": "\\label{summary} We have identified molecular cloud and diffuse interstellar medium material  along the line of sight to the Galactic Center using ISO-SWS data for 3 sightlines.  The amount of molecular cloud material varies across the GC field, and the amount of diffuse ISM material shows little variation.  This is supported by our observations of the 3 \\micron\\ \\water-ice feature which ranges between $\\tau_{3.0}=0.16-0.50$, and the 3.4 \\micron\\ aliphatic hydrocarbon feature which ranges between $\\tau_{3.4}=0.15-0.21$.  The line of sight towards \\sgra\\ contains the largest column of molecular cloud material, relative to GCS3 and 4, apparent from the depth of the 3 \\micron\\ ice feature.  Towards all GC sources discussed in the present paper, CO$_2$ ice is detected, and CO is present mainly in the gas phase.   Towards \\sgra,  solid \\methane\\ has also been firmly identified.  As the amount of solid CO$_2$ relative to \\water\\ is higher toward the Quintuplet sources compared to \\sgra, there may be more than one molecular cloud associated with the overall GC region.  However, the processes that control the solid CO$_2$ abundance are not well understood \\nocite{whittet_etal98} (e.g., Whittet \\etal\\ 1998; Gerakines \\etal\\ 1999)  In many ways, the molecular clouds along the line of sight toward the Galactic Center are similar to local molecular clouds in terms of abundances of key solid state molecules (Table~6).  Water-ice remains the most abundant solid state molecule; the solid CO$_2$/\\water\\ ratio is within the range found for local molecular clouds (Gerakines \\etal\\ 1999); \\methanol/\\water\\ is similar to that observed in quiescent clouds such as Taurus \\nocite{chiar_adamson_whittet96}(Chiar \\etal\\ 1996), and does not resemble the higher abundance observed toward heavily embedded massive protostars.  While solid CO is not detected in the sightlines discussed here, a weak 4.67 \\micron\\ feature has been detected previously (McFadzean \\etal\\ 1989).  The CO/\\water\\ ratio implied is $<10\\%$.  The 3.0 \\micron\\ ice profile is indicative of the temperature, as well as the content of the ices.  The profile  observed toward the GC is distinct from that seen in local molecular clouds such as Taurus.  It is  consistent with cold ($\\sim15$\\,K) ices, similar to local quiescent dense clouds, but   setting this feature apart from those studied in local clouds is substructure at 2.95 \\micron\\ indicative of a relatively high abundance (up to $\\sim30\\%$, relative to \\water) of  \\ammonia\\ in the ices.  Toward local molecular cloud sources, the solid \\ammonia/\\water\\ abundance is observed to be  $\\sim10\\%$ \\nocite{whittet_etal96hh,lacy_etal98}(e.g., Whittet \\etal\\ 1996, Lacy \\etal\\ 1998).  The high ammonia abundance in GC molecular clouds could reflect the higher N/O ratio in the inner Galaxy.  Aliphatic hydrocarbons, first detected via their 3.4 \\micron\\ CH stretch feature, are a widespread component of the diffuse interstellar medium. ISO-SWS has for the first time detected the corresponding C-H deformation modes at 6.85 and 7.25 \\micron\\ in the line of sight towards the Galactic Center.  The ratio of integrated intensities for these three absorption features are consistent with  the predicted values for saturated aliphatic hydrocarbons \\nocite{wexler67} (Wexler 1967).   Hydrogenated amorphous carbon containing a substantial fraction of saturated aliphatic hydrocarbon groups provides a convincing match to the observed hydrocarbon absorption features, and is the likely carrier of these features.  Aromatic hydrocarbons, detected via their 3.3 \\micron\\ C-H and 6.2 \\micron\\ C-C stretching vibrations, do not appear to be evenly distributed throughout the Galactic Center.    Absorption at 3.28 \\micron\\ is detected in the line-of-sight toward GCS3, but not toward \\sgra.  In addition, the 6.2 \\micron\\ absorption feature toward \\sgra\\ is much weaker than that toward the Quintuplet sources.  Although our results suggest that the 3.28 and 6.2 \\micron\\ features are due to the same carrier, it remains unclear whether the C-C stretch absorption feature could instead be produced by dusty C-rich WR stars, and is thus independent of the C-H stretch feature.     \\nocite{chiar_etal98yso} \\nocite{schutte_etal99} \\nocite{tielens_etal91} \\nocite{dartois_etal99} \\nocite{moneti_cernicharo00}"
    },
    "0002/astro-ph0002366_arXiv.txt": {
        "abstract": "The imaging Compton telescope COMPTEL aboard NASA's Compton Gamma-Ray Observatory has opened the MeV gamma-ray band as a new window to astronomy. COMPTEL provided the first complete all-sky survey in the energy range 0.75 to 30 MeV. The catalogue, presented here, is largely restricted to published results. It contains firm as well as marginal detections of continuum and line emitting sources and presents upper limits for various types of objects. The numbers of the most significant detections are 32 for steady sources and 31 for gamma-ray bursters. Among the continuum sources, detected so far, are spin-down pulsars, stellar black-hole candidates, supernova remnants, interstellar clouds, nuclei of active galaxies, gamma-ray bursters, and the Sun during solar flares. Line detections have been made in the light of the 1.809 MeV $^{26}$Al line, the 1.157 MeV $^{44}$Ti line, the 847 and 1238 keV $^{56}$Co lines, and the neutron capture line at 2.223 MeV. For the identification of galactic sources, a modelling of the diffuse galactic emission is essential. Such a modelling at this time does not yet exist at the required degree of accuracy. Therefore, a second COMPTEL source catalogue will be produced after a detailed and accurate modelling of the diffuse interstellar emission has become possible. ",
        "introduction": "COMPTEL has demonstrated that the sky is rich in phenomena that can be studied at MeV energies. A variety of gamma-ray emitting objects are visible either in continuum or line emission. Among the continuum sources are spin-down pulsars, stellar black-hole candidates, supernova remnants, interstellar clouds, nuclei of active galaxies, gamma-ray bursters, and the Sun during solar flares. Line detections have been made in the light of the 1.809 MeV $^{26}$Al line, the 1.157 MeV $^{44}$Ti line, the 847 and 1238 keV $^{56}$Co lines, and the neutron capture line at 2.223 MeV.  COMPTEL has also measured the diffuse interstellar and cosmic gamma radiation, whose properties are described elsewhere (\\cite{strong96,strong99,bloemen99a,kappadath96,weidens99}). For the identification of galactic sources a modelling of the diffuse galactic emission is essential. Such a modelling at this time does not yet exist at the required degree of accuracy.  This paper is restricted to all those sources that have been definitely or marginally detected so far, and provides upper limits to the MeV flux from different types of objects. A second COMPTEL source catalogue will be produced after a detailed and accurate modelling of the diffuse interstellar emission has become possible.  The Compton Observatory was launched on April 5, 1991 by the Space Shuttle Atlantis. During Phase-I of the Compton Observatory Program, a full-sky survey - the first ever in gamma-ray astronomy - was performed. Phase-I ended on November 17, 1992. Observations during the subsequent phases of the program resulted in deeper exposures and complemented the survey. This source catalogue is mainly restricted to the results from the first five years of the mission (up to Phase IV/Cycle-5). In a few cases, more recent results have been added. ",
        "conclusions": "COMPTEL has opened the MeV gamma-ray band (0.75 to 30 MeV) as a new window to astronomy. The data from five years of COMPTEL observations provided the first COMPTEL source catalogue. It is largely restricted to previously published results, and contains firm as well as marginal detections of continuum and line emitting sources, and presents upper limits for various types of objects. The number of most significant detections ($>$ 3$\\sigma$) are 32 for steady sources and 31 for gamma-ray bursters (see Table 13). Six of the listed sources extend over larger areas. Their extent may actually be due to a larger number of so far unresolved point sources. This may be especially true for the Cygnus region in 1.809 MeV, for GRO J1823-12 and for the two HVC complexes. A second COMPTEL source catalogue will be produced after a more accurate modelling of the diffuse interstellar emission has become possible."
    },
    "0002/astro-ph0002199_arXiv.txt": {
        "abstract": "We present the results of a deep $K$-band imaging study which reveals the host galaxies around a sample of luminous radio-quiet quasars. The $K$-band images, obtained at UKIRT, are of sufficient quality to allow accurate modelling of the underlying host galaxy. Initially, the basic structure of the hosts is revealed using a modified Clean deconvolution routine optimised for this analysis. 2 of the 14 quasars are shown to have host galaxies with violently disturbed morphologies which cannot be modelled by smooth elliptical profiles. For the remainder of our sample, 2D models of the host and nuclear component are fitted to the images using the $\\chi^{2}$ statistic to determine goodness of fit. Host galaxies are detected around all of the quasars. The reliability of the modelling is extensively tested, and we find the host luminosity to be well constrained for 9 quasars. The derived average $K$-band absolute $K$-corrected host galaxy magnitude for these luminous radio-quiet quasars is $\\langle M_K \\rangle=-25.15\\pm0.04$, slightly more luminous than an $L^*$ galaxy. The spread of derived host galaxy luminosities is small, although the spread of nuclear-to-host ratios is not. These host luminosities are shown to be comparable to those derived from samples of quasars of lower total luminosity and we conclude that there is no correlation between host and nuclear luminosity for these quasars. Nuclear-to-host ratios break the lower limit previously suggested from studies of lower nuclear luminosity quasars and Seyfert galaxies. Morphologies are less certain but, on the scales probed by these images, some hosts appear to be dominated by spheroids but others appear to have disk-dominated profiles. ",
        "introduction": "Models of the cosmological evolution of quasars often use galaxy mergers as the primary mechanism for quasar activation and require the mass of the structure within which a quasar is formed as a basic parameter \\cite{efrees,haehnelt,percival}. One step towards testing hypotheses about quasar initiation is to answer the question: Is a quasar's luminosity correlated with the luminosity of the structure within which it formed? Such a correlation has been shown to exist for low redshift ($0<z<0.3$) Seyferts and quasars with luminosities $M_V\\simgt-25$ in that there appears to be a lower limit to the host luminosity which increases with quasar luminosity \\cite{mcleod95a,mcleod99}. However this limit is poorly defined, particularly for high luminosity quasars when the strong nuclear component makes it increasingly more difficult to find low luminosity hosts.  Recent work has shown that the majority of nearby galaxies have massive dark objects in their cores, which are suggested to be super-massive black holes potentially capable of powering AGN \\cite{kormendy,magorrian}. These studies have also found evidence for a correlation between the mass of the compact object and the luminosity of the spheroidal component of the host. Assuming a link between nuclear luminosity and black hole mass, the average nuclear luminosity emitted by low redshift quasars is expected to increase with host spheroidal luminosity. In light of this prediction there has been a resurgence of interest in host galaxy studies and recent work \\cite{mclure} has found weak evidence for a correlation in accord with the relations of Magorrian \\etal\\ \\shortcite{magorrian}. However, this correlation relies on spheroid/disk decomposition for two quasars with low nuclear luminosities, and only a small number of luminous radio-quiet quasars were observed.  Host galaxy properties of AGN are known to be correlated with the radio power: radio galaxies tend to be large spheroidal galaxies, while disk galaxies tend to be radio-quiet. Recent evidence suggests that the hosts of radio-loud quasars are also predominantly massive spheroidal galaxies regardless of the nuclear luminosity \\cite{dunlop93,taylor,mclure}. However, studies of radio-quiet quasars with luminosities $M_V\\simgt-25$ have shown that the hosts can be dominated by either disk-like or spheroidal components or can be complex systems of gravitationally interacting components \\cite{taylor,bahcall,boyce,mclure}. There is therefore strong justification for studies to see if these luminosity and morphological trends extend to the hosts of more luminous ($M_V\\simlt-25$) radio-quiet quasars.  There have been many recent detections of host galaxies in the optical thanks to results from HST \\cite{hooper,bahcall,boyce,mclure}, which add to our knowledge from ground-based studies \\cite{malkan84a,malkan84b,hutchings92,veron-cetty}. However, quasar hosts often appear significantly disturbed, as if by interaction or merger which can lead to strong bursts of star-formation and significant extended line and blue continuum emission at optical wavelengths which are not indicative of the mass of the underlying host. The nuclear-to-host light ratio in the optical is also typically higher than at longer wavebands.  These problems can be circumvented by observing in the infrared where the contrast between host and nuclear component is improved and the emission associated with star bursting activity is largely absent: the $K$ magnitude is a better measurement of the long-lived stellar populations in the host \\cite{bruzual}. Previous observations in the infra-red have been successfully used to determine quasar host galaxy luminosities and morphologies (McLeod \\& Rieke 1994a;b; Dunlop \\etal\\ 1993; Taylor \\etal\\ 1996). However, recent advances in telescope design, in particular the advent of adaptive optics systems such as the tip-tilt system on UKIRT produce clearer images of quasars and enable accurate point spread functions (psfs) to be more readily obtained as differences between successive observations are reduced. Such advances coupled with improved analysis techniques mean we are now able to reveal the host galaxies of luminous quasars with $M_V\\simlt-25$ using infra-red observations.  In order to obtain enough luminous radio-quiet quasars our sample was forced to cover redshifts $0.26\\le z\\le0.46$. At such redshifts, with typical seeing, the structure of the host galaxy is hidden in the wings of the psf from the nuclear component. There are two main ways of proceeding: either the psf can be deconvolved from the quasar light to directly observe the host galaxy, or known galaxy profiles can be used to model the hosts, a nuclear component can be added in and the profiles can be fitted to the data. Because the host galaxies sometimes have disturbed morphologies indicative of violent mergers, it is difficult to assume a form for the galaxy. However without such modelling, it is not easy to determine the contribution of the host to the light from the centre of the quasar and deconvolution routines tend to produce biased solutions which may alter important features.  For the analysis of our quasar sample, an approach is adopted which uses both methods. Initially (Section~\\ref{sec:simple}) the images were restored using a deconvolution algorithm, based on the Clean algorithm \\cite{hogbom}, developed for this problem, which will be described elsewhere \\cite{percival_clean}. This routine was used to reveal the extent to which the `nebulosity' around the point source is disturbed. Deconvolution of the light from two of our quasars reveals violently disturbed host galaxies indicative of close merger events. In the remainder of our sample, the non-nuclear light is more uniformly distributed around the centre of the quasar. We should note that the resolution provided by this deconvolution technique is probably not sufficient to reveal evidence for weak mergers, where the host galaxy is only slightly disturbed.  Where the image-restoration routine revealed approximate elliptical symmetry in the non-nuclear component, 2D galaxy profiles were fitted to the hosts. Analysis of non-interacting, low redshift galaxies has shown that an empirical fit to both disk and spheroidal systems is given by: \\begin{equation} \\mu=\\mu_o\\exp\\left[-\\left(\\frac{r}{r_o}\\right)^{1/\\beta}\\right]. \\label{eq:galprofile} \\end{equation} where $\\mu$ is the average surface brightness in concentric elliptical annuli around the core, and $r$ is the geometric average of the semi-major and semi-minor axes.  Model images were carefully created using this profile and were tested against the data using the $\\chi^2$ statistic to determine goodness of fit. Five host parameters were required, the half-light radius, integrated luminosity, axial ratio, angle on the sky, and the power-law parameter of the galaxy $\\beta$, as well as the nuclear-to-host ratio. Section~\\ref{sec:model} describes the modelling procedure in detail, and in Section~\\ref{sec:results} the best-fit parameters are presented for the host galaxies.  Much previous work has produced ambiguous results because of a lack of error analysis and insufficient testing of the modelling. A detailed analysis of the reliability of the 2D modelling method used in this paper has therefore been undertaken and is presented in Section~\\ref{sec:test}. Although hosts are detected in all of our sample, the upper limit of the host luminosity is only usefully constrained for 9 of the 12 quasars modelled (see Section~\\ref{sec:calc_lum}). Similar analysis of the best-fit $\\beta$ parameter which determines the morphology of the host reveals that this parameter is, unsurprisingly, more poorly constrained than the luminosity. However, we have created Monte-Carlo simulations of images with the same signal-to-noise as the original images (Sections~\\ref{sec:mock}~\\&~\\ref{sec:mock2}). By analysing these images using exactly the same procedure as for the original data we find that it is possible to distinguish between disk and spheroidal structure.  Unless stated otherwise we have adopted a flat, $\\Lambda=0$ cosmological model with $H_0=50$km\\,s$^{-1}$Mpc$^{-1}$ and have converted previously published data to this cosmology for ease of comparison. ",
        "conclusions": "We have presented the results from a deep $K$-band imaging study designed to reveal the host galaxies of quasars with higher luminosities than targeted by previous studies. Extending host-galaxy studies to these quasars was made possible by the stability provided by the tip-tilt adaptive optics system at UKIRT, which enabled accurate psf measurements to be obtained for the deep quasar images. We have been able to resolve host galaxies for all of our sample.  The principle conclusion of this study is that the luminous quasars in this sample have host galaxies with similar luminosities to quasars of lower total luminosity. Derived nuclear-to-host ratios are therefore larger than those obtained by previous work, and place these quasars beyond the upper limit suggested by studies of quasars with lower total luminosities. Host morphologies are less certain, but there is weak evidence that the hosts of these quasars can be dominated by either disk-like or spheroidal profiles on the scales probed by these images."
    },
    "0002/astro-ph0002220_arXiv.txt": {
        "abstract": "We present new observations of two anomalous X-ray pulsars (AXP) \\axpa\\ and \\axpb\\ made in 1998 with the \\asca. The energy spectra of these two AXPs are found to consist of two components, a power-law and a blackbody emission from the neutron star surface. These observations, when compared to earlier \\asca\\ observations in 1994 show remarkable stability in the intensity, spectral shape and pulse profile. However, we find that the spin-down rate in \\axpb\\ is not constant. In this source, we have clearly identified three epochs with spin-down rates different from each other and the average value. This has very strong implications for the magnetar hypothesis of AXPs. We also note that the spin-down rate and its variations in \\axpb\\ are much larger than what can normally be produced by an accretion disk with very low mass accretion rate corresponding to its low X-ray luminosity. ",
        "introduction": "Some X-ray pulsars are known to have remarkable similarity in their properties which are different from other binary or isolated X-ray pulsars (\\cite{mere95b}). The properties common to most of these objects are a) pulse period in a small range of 5--12 s, b) monotonous spin down with ${\\rm P} / \\dot {\\rm P}$ in the range of $5\\times10^{11} - 1.3\\times10^{13}$ s c) identical X-ray spectrum consisting of steep power-law ($\\Gamma = 3-4$) and black body component (kT $\\sim 0.5$ keV), d) stable X-ray luminosity (10$^{34}-10^{36}$ ergs s$^{-1}$) for years, e) faint or unidentified optical counterpart, and f) no evidence of orbital motion. The sources also have a galactic distribution, most of these are within $|{b}|\\leq 0.5^\\circ$ and all are probably young ($\\sim 10^4$ yr) because of their association with SNR or molecular clouds. The objects in which all the properties mentioned above have been observed are \\axpa, 1E~2259+586, \\axpb, 1RXS~J170849.0$-$400910 and 1E~1841$-$045 (Kes~73). Two more objects, AX~J1845.0$-$0300 (\\cite{tori98}) and RX~J0720.4$-$3125 (\\cite{habe96}), also probably belong to the same class but to establish their AXP candidacy, more X-ray observations are required to measure their pulse period variations, search for possible pulse arrival time delay and investigate the flux stability. Classification of an object as AXP only from some properties similar to the above is not very firm. 4U~1626$-$67, probably a binary system, showed both spin-up and spin-down (\\cite{chak97}) and also has an optically bright accretion disk (\\cite{midd81}). Another one object RX~J1838.4$-$0301 (\\cite{schw94}), does not have stable intensity and pulsations are also not always detectable (\\cite{song99}). Therefore, these two objects are not AXPs.  Considering the strong similarity between these handful of sources, it has been proposed that they have same physical nature and different scenarios have been proposed to explain the observed properties. The prominent models are a) accretion from low mass binary companion (\\cite{mere95b}), b) single neutron star accreting from molecular cloud, or a product of common envelope evolution (Thorne-$\\rm \\dot{Z}$ytkov object) of close high mass X-ray binaries in which a solitary neutron star accretes matter from a fossil disk (van Paradijs, Taam, \\& van den Heuvel \\cite{vanp95a}; Ghosh, Angelini, \\& White \\cite{ghos97}), and c) extremely high magnetic field neutron star radiating X-rays due to magnetic field decay (\\cite{thom96}). Unlike the radio pulsars and rotationally powered X-ray pulsars, in the AXPs, the spin-down rate is not large enough to power the observed X-ray emission.  Among the 90 or so known X-ray pulsars (\\cite{naga99}), direct evidence of binary nature is known for more than 35 sources (\\cite{vanp95b}). Including the indirect evidences this number can be upto about 65 and 7 pulsars are isolated stars in SNR and are powered by rotational energy losses. In the rest of the pulsars, in which no binary signature is known, it is often due to lack of sufficient observation. However, the 7 objects which are either AXPs or candidate AXPs (or 10 if we include the 3 Soft Gamma-ray Repeaters in which pulsations have been detected), no binary signature has been found in spite of extensive searches. The strong upper limit on pulse arrival time delay that has been obtained in some of these sources strongly suggests non-binary nature for the AXPs. In addition, the spin change behaviour of the AXPs is also remarkably different from accreting pulsars (\\cite{bild97}). In almost all the accreting X-ray pulsars, both spin-up and spin-down episodes have been seen which may be randomly distributed (in persistent sources) or spin-downs in quiescence followed by rapid spin-ups during bright transient phases (in transient pulsars) or long monotonic spin-up and spin-down episodes accompanied by spectral and luminosity changes (e.g. 4U~1626$-$67, \\cite{yiiv99}).  The AXPs are in many respect also similar to the X-ray counterparts of the Soft Gamma-ray Repeaters (SGR). The X-ray spectral and timing properties of these two type of objects have strong similarities, half of the AXPs and SGRs are associated with supernova remnants. This has lead to the suggestion that the AXPs are also magnetars in which the X-ray emission is due to magnetic field decay (\\cite{thom96}). The main difference between these two type of objects is the non detection of SGR bursts from the AXPs. However, considering the rarity of the SGR activity among the established SGRs (\\cite{kouv96}), the absence of bursts from AXPs is not a serious issue. From a relatively young age of the AXP, 1E~1841$-$045 in the supernova remnant Kes~73, Gotthelf, Vasisht, \\& Dotani (\\cite{gott99}) proposed that in the evolutionary track, the AXPs are an early quiescent state of the SGRs. Stability of the X-ray emission properties (spin-down rate, luminosity, spectral shape and pulse shape and fraction) is usually mentioned as one important aspect of the AXPs though one has to compare between observations made with different instruments for which the energy band, energy resolution and sensitivity are not identical. To make a rigorous comparison in the stability of the X-ray emission properties we have made new observations of two AXPs with the \\asca, four years after two previous observations reported by \\cite{whit96} and Corbet \\& Mihara (\\cite{corb97}). The aim was to critically examine the stability of the X-ray emission pattern and more pulse period measurements which may provide support to either the accretion powered or the magnetar hypothesis for these objects.  The source \\axpa\\ is close to a long period binary pulsar RX J0146.9+6121 and in the \\exo\\ observations of this field in which pulsations were first discovered, pulsations from both the sources were observed simultaneously (\\cite{isra94}). With \\exo, the 8.7 s pulsations in this source were detected only in the 1.0--4.0 keV range. A binary nature of the system was preferred in spite of a large X-ray to optical flux ratio and absence of pulse arrival time delay. The source was subsequently observed with \\ros\\ and the pulse period was found to be very close to the \\exo\\ measurement (\\cite{hell94}). \\asca\\ observation in 1994 confirmed the steady spin-down trend and defined the spectral character clearly. A model consisting a 0.4 keV blackbody and a power law with a photon index of 3.7 was found to describe the spectrum well. From a small radius of a few km of the black body emission zone (which probably is on the surface of the neutron star) \\cite{whit96} suggested that the black-body component is more likely to be due to a spherical accretion rather than accretion from a disk. A small pulse fraction and energy dependent pulse profile, double peak at low energy and single peak at high energy is characteristic of this pulsar. From \\xte\\ observations, the upper limit on the pulse arrival time delay was determined to be 260 ms in the 70 s to 2.5 days range thereby ruling out all types of binary companions except white dwarf or low mass He main sequence star (\\cite{wils99}). \\sax\\ observations during 1997--98 confirmed the spectral characteristics, pulse profiles and spin-down trend (\\cite{isra99b}).  Observations spanning 20 years during 1979--1998, with the \\ein, \\exo, \\ros, \\asca, \\sax\\ and \\xte\\ show an overall constant spin-down trend with $\\dot {\\rm P}  = 2.2 \\times 10^{-12}$ s s$^{-1}$.  Pulsations in the source \\axpb\\ were discovered from observations with the \\ein\\ observatory in 1979 and was confirmed by \\exo\\ observation in 1985 (\\cite{sewa86}). The energy spectrum was found to be a power-law type ($\\Gamma$ = 2.26) with low energy absorption (N$_{\\rm H} = 1.6 \\times 10^{22}$ atoms cm$^{-2}$). The relatively harder power-law spectrum and a candidate optical counterpart lead to the speculation that this can be a Be star binary. Several \\gin\\ observations established a secular spin-down trend with $\\dot {\\rm P} = 1.5 \\times 10^{-11}$ s s$^{-1}$ similar to a few other soft spectrum pulsars (\\cite{corb90}). Subsequent \\ros\\ observations however revealed an increase in the spin-down rate which is remarkably different from other established AXPs (\\cite{mere95a}). The low energy part of the spectrum was first accurately measured with the \\asca\\ in 1994 (Corbet \\& Mihara \\cite{corb97}). However, it could not distinguish between power-law ($\\Gamma = 3.34$) and a combination of power-law ($\\Gamma$ = 2.0) and black-body (kT = 0.55 keV). During the \\ros\\ and the \\asca\\ observations, the spin-down rate remained at a higher level of $3.3 \\times 10^{-11}$ s s$^{-1}$. A decrease in intensity by a factor of 3 compared to the \\exo\\ observation was also noticed. \\sax\\ observation in 1997 showed that a combined black-body and power-law model fits the data well. The size of the black-body emission zone was found to be of the order of km$^2$, identical to other AXPs (\\cite{oost98}). Compared to other AXPs, the pulse fraction was found to be much larger ($\\sim$ 65\\%) in \\axpb\\ and has little energy dependence. The power-law component is also relatively harder ($\\Gamma = 2 \\sim 2.5$) compared to the other AXPs ($\\Gamma = 3 \\sim 4$). \\xte\\ observations in 1996--97 showed yet another change in the spin down rate, now close to that in 1980's. Long \\xte\\ observations established a small upper limit of 60 ms for any pulse arrival time delay for orbital period in the range of 200 s to 1.5 days. Based on this, any binary companion other than low mass helium burning star in a face on system has been ruled out (\\cite{mere98}). ",
        "conclusions": "Using multiple observations of two AXPs with the {\\it \\asca} we have found remarkable stability in the intensity, spectral shape and pulse profile over a 4 years period. For the source \\axpb, we have confirmed that similar to other AXPs, the spectrum consists of a power-law component and a black body component. The spin-down trend of \\axpa\\ is consistent with a constant rate whereas in \\axpb\\ we have clearly identified three different spin-down epochs. We have shown that the fast spin-down of some of the AXPs is difficult to achieve with disk accretion. Hence the common envelope evolution scenario of AXPs is unlikely to be the case, unless the spin-down is due to wind outflow from the disk. In this case also, the stability of X-ray emission in spite of varying spin-down rate remains unexplained. The significantly different spin-down rates of \\axpb\\ in different epochs are also difficult to reconcile in the magnetar model, unless precession is at work. But the flux stability and lack of variation of the spectral parameters appear to favour the magnetar model."
    },
    "0002/astro-ph0002016_arXiv.txt": {
        "abstract": "Recent HST observations suggest that the NLR in Seyfert Galaxies can be the result of interaction between jet and external inhomogeneous medium; following this suggestion we perform numerical simulations considering the impact of a radiative jet on a dense cloud. We approach the problem adopting  a hydrodynamical code, that consents us to study in detail the jet hydrodynamics, while we choose a more simplified treatment of radiative  processes, in order to give a qualitatively good interpretation of the emission processes.  Our three main purposes are: i) to reproduce in our simulations the physical conditions observed in the NLR of Seyfert Galaxies, ii) to obtain physical constraints of the jet parameters and iii) to study the jet capacity to photoionize the surrounding medium.  We find that the jet-cloud interaction leads to clumps of matter with density, temperature and velocity that agree with observations. Conversely, the photoionizing flux radiated by the jet-induced shocks does not appear to be sufficient to account for the NLR line luminosity but it may produce local and transient effects on the NLR ionization balance.  Finally, the observational requirements can be matched only if jets in Seyfert galaxies are relatively heavy, $\\rho_{\\rm jet}\\, \\gapp 1 $ cm $^{-3}$, and with velocities smaller than $\\lapp \\, 50,000\\,$km s$^{-1}$, very different from their counterparts in radio-galaxies. ",
        "introduction": "Extensive HST emission-line imaging of Seyfert galaxies has for the first time resolved details of the structure of their Narrow Line Regions (NLR). In several cases cone-like morphologies have been revealed, similar in shape to - but of much smaller linear extent than - the Extended Narrow Line Regions (ENLR) seen in the lower resolution ground based images (Wilson \\& Tsvetanov, 1994 and references therein). In the standard model of the NLR, the UV emission of the nucleus is responsible of photoionizing the Interstellar Medium (ISM) of the host galaxy. These conical distributions of the ionized gas have been interpreted as a confirmation of the anisotropy of the nuclear radiation field which, in the framework of the unified scheme for Seyfert galaxies (e.g. Antonucci 1993) is caused by the shadowing of an obscuring circumnuclear torus. However, in galaxies with linear radio structures, the morphology of the emission-line region appears to be directly related to that of the radio emission. In particular, in Seyferts with radio jets (e.g. Mrk 3, Mrk 348, Mrk 6, Mrk 1066, ES0 428-G14), the NLR itself appears jet-like and is spatially coincident with the radio jet, while the emission-line region takes a different form when a radio lobe is present (e.g. Mrk 573, Mrk 78, NGC 3393): arc-like shells of emission, very reminiscent of bow-shocks, surround  the leading edge of the lobes (Capetti et al. 1995a, 1995b, 1996; Falcke et al. 1996, 1998). This dichotomy in radio and emission-line morphology is reflected in their different scales: bow-shock structures cover several kiloparsecs, while the jet-like features extend only over a few hundred parsecs. The simplest interpretation of this radio-to-optical correspondence is that the radio emitting outflow creates an expanding and cooling gas halo. The compression induced by the outflow causes the line emission to be highly enhanced in the regions where the jet-cloud interactions occur. A clear confirmation of this scenario came recently from HST spectroscopy of Mrk 3 (Capetti et al. 1999): its NLR has velocity field characteristic of a cylindrical shell expanding at a rate of 1700 km/s. They interpreted this as the consequence of the rapid expansion of a hot gas cocoon surrounding the radio-jet, which compresses and accelerates the ambient gas.  HST observations also provided evidence for spatial variations in the NLR ionization structure. In NGC 1068 the material located along the radio jet is in a much higher ionization state than its surroundings. This might suggest the presence of a local source of ionization which dominates over the nuclear radiation field (Capetti, Axon and Macchetto 1997; Axon et al. 1998). In other sources, too distant for such a detailed analysis, the radial variations of the ionization parameter are generally much flatter than expected from pure nuclear photoionization on the basis of the measured density gradients (Capetti et al. 1996, Allen et al. 1999) requiring again a local source of ionizing photons. An appealing possibility of interpreting these data is to invoke the ionizing effects of shocks, originated by  jet-cloud interactions: if these shocks are  fast enough (velocities $>$ a few hundred km s$^{-1}$) the hot, shocked gas could produce a significant flux of ionizing photons (Sutherland, Bicknell and Dopita 1993; Dopita and Sutherland 1995, 1996). Direct evidence for this emission has been found by Axon, Capetti and Macchetto (1999) who showed that the radio-jets in the Seyfert 2 galaxies Mrk 348 and Mrk 3 are associated with an extended linear structure in UV and optical continua. In this picture, the radio-jet would not only determine the morphology of the NLR but is physically involved in its ionization. A radio imaging survey of the CfA sample of Seyfert (Kukula et al. 1995) shows that radio linear structures are present in a large fraction of sources (more than 50\\%) suggesting that such an interaction is likely to be a quite common phenomenon in this class of objects.  The jet interaction with the external medium is clearly a complex physical problem which involves both a hydrodynamical study of the jet propagation as well as a detailed understanding of the microphysics of the induced shocks, which might also be magnetized, and of the radiative processes.  In the framework of Seyfert galaxies this issue has been tackled by several authors (Dopita and Sutherland 1995, 1996, Evans et al. 1999, Wilson and Raymond 1999, Allen et al. 1999). Their focus is however mainly on the shocks properties with a very detailed treatment of the emission mechanisms, with simplifying assumptions about the hydrodynamics (e.g. plane parallel geometry, steady-state shock). The comparison with the observations is based on the emitted spectrum and in particular on diagnostic line-ratios, particularly with the aim of distinguishing the different signatures of nuclear versus local photoionization.  In this paper we follow a complementary, albeit different, approach by studying in detail the jet hydrodynamics, while adopting a simplified treatment of the radiative processes, as we employ an equilibrium cooling function in an optically thin approximation. This approach allows us to compare the results of simulations with the observed properties of NLR, in particular their morphology, the expansion velocities and the characteristic values of gas density and temperature. More precisely we consider the interaction of the jet with an inhomogeneity in the external medium (cloud) and our aim is that of constraining the jet and cloud physical parameters for which it is possible to reproduce the observed conditions. In this way, in addition of getting a better understanding of the NLR physics, we can also obtain information on the jet properties from the NLR data. Moreover, we can  calculate the fraction of the jet power converted in radiation by shocks, resulting from the interaction of the jet with the environment. We then get from the global dynamics a conversion efficiency from kinetic to radiative power and we can determine whether the jet itself, via shocks, can provide an {\\it in situ} photoionization source  for the NLR emitting material, as discussed above.  Steffen et al. (1997a) have used a rather similar approach with the main difference that they considered the jet propagating into a uniform medium. It seems that in this situation it is impossible to reach the high densities typical of the NLR with jet-like emission (see discussion below) on which we will focus in the present paper. This is because, at low density, radiation is not efficient enough to give the needed compression factors. The conditions of the emitting material obtained by Steffen et al. seem to be appropriated for the case of the more extended (lobe-like) line emission structure.  Steffen et al. (1997b) considered also jet-cloud interactions mainly from an analytical point of view. They found that when a jet interacts with a large number of clouds the most relevant effects on the NLR structure are due to the most massive clouds located along the jet path. This lead us to our choice for the geometry of the simulation in which the jet hits a single dense cloud.  The paper is structured as follows: In the next section (Sect. 2), we describe the basic physical problem and the observational constraints, while the equations used and the method of solution are examined in Sect. 3 and 4; the results of simulations are discussed in Sect. 5; conclusions are drawn in Section 6. ",
        "conclusions": "We have studied in detail the dynamics of the interaction of a jet with a large cloud pre-existing in the ISM in order to find the conditions for which it is possible to reproduce the main physical parameters of the NLR emitting material.  Following the suggestion by Steffen et al. (1997b) that the most relevant effects of the interaction arise when a jet hits dense massive clouds, we adopted a quite simplified geometry of a single gas condensation with can be astrophysically identified with a giant molecular cloud. As the interaction last for a time considerably longer than the cloud crossing time more than one cloud will be interacting at any given time and they will display simultaneously the different evolutionary stages of the interaction. Furthermore the characteristic jet-line structure of the NLR is thus reproduced. In any event, this case, i.e. the head-on collision with a large cloud, is the most efficient case of interaction, for the compression, acceleration and heating of the NLR material.  We concentrated  our efforts on the exploration of the parameter plane ($v_{\\rm jet}, \\ \\rho_{\\rm cloud})$, since the other parameters, on which the simulation depends, have little influence on the properties of the optically emitting material. We have found that the condition for obtaining values of density, temperature and velocity in the observed range can be translated in a condition on the parameter  $\\tau$, which is the ratio of the cloud crossing timescale to the radiative timescale ($0.3 < \\tau < 0.55$) and which depends on our two fundamental parameters $v_{\\rm jet}$ and  $\\rho_{\\rm cloud}$. For small values of $\\tau$, radiation is inefficient and it is not possible to produce regions dense enough, while, on the other hand, for large values of $\\tau$, the cloud is too dense and the obtained velocities are too low. We have explored a range of cloud densities which can be considered typical of GMCs and, for this range, the jet velocities  span an interval from  $4000\\,$km s$^{-1}$ to $55,000\\,$km s$^{-1}$.  The jet kinetic power corresponding to these combinations of parameters (for a jet density of 1 cm$^{-3}$) ranges from $3.2 \\times 10^{41}$ to $8 \\times 10^{44}$ erg s$^{-1}$, in general agreement with the estimates of Capetti et al. 1999 for Mrk 3. For jet density much lower than 1 cm $^{-3}$, however, in order to match the observed NLR conditions we would need a correspondingly higher velocity and therefore untenable requirements on the kinetic power which grows with $v_{\\rm jet}^3$. We conclude that jets in Seyfert galaxies are unlikely to have densities much lower than 1 cm $^{-3}$ and velocities higher than $50,000\\,$km s$^{-1}$, and therefore they are very different from their counterparts in radio-galaxies in which densities are much lower and velocities are relativistic.  Concerning radio-galaxies we can speculate that with lower jet densities and higher velocities, the gas postshock temperatures and radiative time would be increased with respect to the case of Seyfert galaxies and therefore the conditions for having efficient line emission would be more difficult to meet. In addition the different properties of the jet environment in the elliptical galaxies hosting radio-galaxies render encounters with gas condensations less likely to occur. This probably explain why the association between radio and line emission although often present in radio-galaxies (e.g. Baum and Heckman 1989) is not as strong as in Seyfert galaxies.  Finally, the study of the global dynamics allowed us to have estimates of the overall efficiency of the conversion of kinetic to high frequency radiative power in the shocks that form in the interaction between jet and ambient medium. We have found that the efficiency is increased by the presence of the cloud, its peak value is 0.1 - 2 \\% , its typical value is much lower $\\sim 10^{-4} - 5 \\times 10^{-3}$ and it decreases with the jet power. These results lead us to the conclusion that radiation emitted in shocks can be only a small fraction of the overall ionization budget of the NLR, although it can have local and transient important effects."
    },
    "0002/astro-ph0002308_arXiv.txt": {
        "abstract": "We compare the structure and substructure of dark matter halos in model universes dominated by collisional, strongly self interacting dark matter (SIDM) and collisionless, weakly interacting dark matter (CDM). While SIDM virialised halos are more nearly spherical than CDM halos, they can be rotationally flattened by as much as 20\\% in their inner regions.  Substructure halos suffer ram-pressure truncation and drag which are more rapid and severe than their gravitational counterparts tidal stripping and dynamical friction. Lensing constraints on the size of galactic halos in clusters are a factor of two smaller than predicted by gravitational stripping, and the recent detection of tidal streams of stars escaping from the satellite galaxy Carina suggests that its tidal radius is close to its optical radius of a few hundred parsecs --- an order of magnitude smaller than predicted by CDM models but consistent with SIDM.  The orbits of SIDM satellites suffer significant velocity bias $\\sigma_{_{\\rm SIDM}}/\\sigma_{_{\\rm CDM}}=0.85$ and are more circular than CDM, $\\beta_{_{\\rm SIDM}} \\approx 0.5$, in agreement with the inferred orbits of the Galaxy's satellites. In the limit of a short mean free path, SIDM halos have singular isothermal density profiles, thus in its simplest incarnation SIDM is inconsistent with galactic rotation curves. ",
        "introduction": "The nature of dark matter is still far from being resolved. Primordial nucleosynthesis and observational data suggest that the baryonic material accounts for just a fraction of the matter density in the universe.  Fundamental particles remain the most likely candidate for the dark matter and much effort has been devoted to researching a class of weakly interacting, collisionless dark matter (CDM) (\\eg, Davis \\etal{} 1985). However, the hierarchical gravitational collapse of cold collisionless particles leads to dense, singular dark matter halos -- a result that is central to several fundamental problems with this model on small scales (\\eg~Hogan \\& Dalcanton 2000 and references within).  It may be possible to solve the current problems with CDM by appealing to extreme astrophysical processes. Alternatively, we can explore other dark matter candidates that behave differently on non-linear scales. One possibility is strongly self interacting dark matter (hereafter SIDM). Originally proposed to suppress small scale power in the standard CDM model (Carlson \\etal{} 1992, Machaceck \\etal{} 1994, de Laix \\etal{} 1995), SIDM was recently revived by Spergel \\& Steinhardt (1999) to solve some of the outstanding problems with CDM.  The behaviour of this component depends on the particles' collisional cross-section. Large cross-sections imply short mean free paths, so that the dark matter can be described as a fluid that does not cool but can shock heat. Particles with a mean free path of order the scale length of a dark matter halo offers the possibility of conductive heat transfer to the halo cores (Spergel \\& Steinhardt 1999).  In this Letter we contrast the dynamics and structure of ``halos within halos'' between collisional and collisionless dark matter and compare predictions with current observational constraints.  \\ ",
        "conclusions": "The properties of dark matter halos of strongly interacting particles are markedly different from their collisionless counterparts. SIDM halos are close to spherical with a modest degree of rotational flattening.  Observations of halo shapes cannot currently distinguish between the models examined here; however, future lensing observations will determine if SIDM is a viable dark matter candidate.  Halos within halos suffer ram-pressure truncation that decreases their sizes to less than the tidal radius. Current observational data on galactic halos in clusters and satellite galaxies in the Galactic halo are naturally reproduced in SIDM models: the extent of Carina's halo is an order of magnitude smaller than predicted by CDM.  Ram-pressure drag creates significant velocity and orbital bias in the substructure halos which sink on a short timescale---of order the crossing time---independent of their mass. Another positive feature of SIDM is the ability to produce satellite systems on near circular orbits which are very rare in CDM models.  Both CDM and SIDM with a large cross-section fail to reproduce observed rotation curves of dwarf and LSB galaxies. We have seen that the final density profiles are sensitive to the shape of the initial fluctuations: more violent collapses end up with constant density cores.  Alternatively, SIDM with a mean free path between kiloparsec and megaparsec scales may solve this problem (Spergel \\& Steinhardt 1999).  In this case, particles could transfer heat to the cold central regions that occur in standard CDM collapses, creating an initial expanding phase with lower central density. It is not obvious that a cold core would be generated and maintained in a hierarchical scenario since the dense mini-halos collapsing at high redshift may form singular isothermal structures. The dense substructure halos would rapidly sink to the centres of the parent halos by hydro-dynamical drag, depositing high density low entropy material and conserving isothermal profiles.  Simulating intermediate mean free paths is relatively straightforward. One technique would be to use the neighbour lists to choose random particles to collide (Burkert 2000).  Simulations in progress will demonstrate whether SIDM can reproduce the observed rotation curves of dwarf galaxies. A solution to this problem will naturally resolve the abundance of dark matter substructure in the Galactic halo since substructure with shallow potentials would be easily disrupted."
    },
    "0002/astro-ph0002078_arXiv.txt": {
        "abstract": "We have obtained 1.64, 1.90 and 2.15 $\\mu$m narrow-band images of five T Tauri stars in the TW Hya Association (TWA) using the Near-Infrared Camera and Multiobject Spectrometer aboard the {\\it Hubble Space Telescope}.   Most of the T Tauri stars in our study show evidence of absorption by H$_2$O vapor in their atmospheres; in addition, the low-mass brown dwarf candidate, TWA 5B, is brighter at 1.9 $\\mu$m than predicted by cool star models that include the effects of H$_2$O vapor but neglect dust.  We conclude that the effect of atmospheric dust on the opacity is important at 1.9 $\\mu$m for TWA 5B, the coolest object in our sample.  The available evidence suggests that the TWA is 5--15 MY old.  Comparison of the colors of TWA 5B with theoretical magnitudes as a function of age and mass then confirms previous claims that TWA 5B is substellar with a mass in the range 0.02--0.03 $\\,M_\\odot$.  The accurate single-epoch astrometry of the relative positions and separation of TWA 5A and TWA 5B reported here should permit the direct measurement of the orbital motion of TWA 5B within only a few years. ",
        "introduction": "For more than a decade, the young star TW Hya has been an enigma since it lies in a region of sky apparently devoid of the raw materials to form stars, nearly 13$^\\circ$ from the nearest dark cloud, yet it is unambiguously a classical T Tauri star (\\cite{ruci1983}) surrounded by a great deal of cold  dust (\\cite{wein1989}) and gas (\\cite{zuck1995,kast1997}). Recently, the TW Hya mystery was solved:  TW Hya, along with other T Tauri stars found in an area of $\\sim$100 square degrees of the southern sky (\\cite{dela1989,greg1992}), compose a uniquely close association of young stars known as the TW Hya Association (\\cite{kast1997}).  At a mean distance of only $\\sim$55 pc, the TW Hya Association (hereafter TWA) is almost three times closer than the next nearest known region of recent star formation. Given the likely age ($\\sim$10 MY) of the TWA, these stars could harbor very young planetary systems with fully formed giant planets or low mass, brown dwarf companions, and may still be surrounding by circumstellar disks.  In fact, there is substantial evidence for circumstellar gas and dust around several of these stars (\\cite{wein1989,zuck1993,zuck1995,kast1997}). The relative proximity and the absence of significant interstellar or intra-molecular cloud extinction in the direction of the TWA make the prospects for detecting substellar companions around these nearest T Tauri stars much better than the prospects for similar searches for young low mass companions around T Tauri stars in Taurus-Auriga, Chamaeleon, Lupus or Ophiuchus, the next closest regions of star formation. In a recent study of the TWA, Webb et al.~(1999) identified a total of at least 17 sources as members of the TWA.  In addition, Lowrance et al.~(1999) and Webb et al.~have reported the discovery of a likely low mass brown dwarf companion (M $\\simeq$ 0.02 $\\,M_\\odot$) to TWA 5A (= CD$-$33$^\\circ$7795) in a combination of ground-based and {\\it Hubble Space Telescope} ({\\it HST}) observations. TWA 5B is found almost 2\\arcsec\\ from TWA 5A; thus, despite its relative physical proximity ($\\sim$100 AU) to the primary, TWA 5B is amenable to spectroscopic and astrometric studies, uncontaminated by light from TWA 5A.  In this paper, we report results from imaging the fields around five stars in the TWA, including the TWA 5 system, using the Near Infrared Camera and Multi-Object Spectrometer (NICMOS) and the {\\it HST}.  The goal of this program was to search for companions around these stars.  Our choice of three narrowband filters centered at 1.64, 1.90 and 2.15 $\\mu$m was designed to enable us to identify cool and low surface gravity objects, including substellar mass companions, through their likely strong signatures of H$_2$O absorption at 1.9 $\\mu$m. ",
        "conclusions": "To the sensitivity limits of these data, our images reveal no detectable circumstellar disks or infrared reflection nebulae, and no low mass stellar or substellar companions around stars in the five studied TWA systems other than the previously discovered TWA 5B. As for TWA 5B itself, our results suggest that this object has a mass in the range of 0.02--0.03 $\\,M_\\odot$, in good agreement with the work of Lowrance et al.~(1999).  Finally, while our single epoch observations cannot demonstrate or measure the orbital motion of TWA 5B, they are more than accurate enough to permit the measurement of this motion, in combination with future epoch {\\it HST} or adaptive optics, ground-based observations, with a baseline of only about a year."
    },
    "0002/astro-ph0002002_arXiv.txt": {
        "abstract": "We apply elliptical isothermal mass models to reproduce the point source properties, i.e. image positions, flux density ratios and time delay ratios, of the quadruple lens B1608+656. A wide set of suitable solutions is found, showing that models that only fit the properties of point sources are under-constrained and can lead to a large uncertainty in the prediction of H$_\\circ$. We present two examples of models predicting H$_\\circ\\!=\\!100{\\rm km\\,s^{\\!-\\!1}Mpc^{\\!-\\!1}}$ ($\\chi^2\\!=\\!4$) and H$_\\circ\\!=\\!69{\\rm km\\,s^{\\!-\\!1}Mpc^{\\!-\\!1}}$ ($\\chi^2\\!=\\!24$). ",
        "introduction": "Relative positions, flux ratios and time delays of the 4 images in B1608+656 have been presented here by Fassnacht (1999) and references therein ({\\it cf} Table~\\ref{obs}). Koopmans \\& Fassnacht (1999) have concluded H$_\\circ\\!=\\!59^{+7}_{-6}{\\rm km\\,s^{\\!-\\!1}Mpc^{\\!-\\!1}}$ within the context of a family of parametrized, isothermal models. Here, we investigate whether a larger set of models allows a wider range of Hubble constants. ",
        "conclusions": "A variety of parametrized models can reproduce the point source properties of B1608+656. This precludes an accurate determination of H$_\\circ$. To break the degeneracy, extra constraints, associated with the extended emission of the source, have to be incorporated. It is also helpful to specify the distribution of dark matter on larger scale than the image distribution. A similar conclusion has been drawn by Williams \\& Saha (1999) using pixellated models. \\begin{figure} \\vspace*{-0.5cm} \\begin{center} \\leavevmode \\epsfysize=1.55in \\epsfbox{gsurpi1.eps} \\vspace*{-0.2cm} \\caption{Surface mass density in Models A and B} \\vspace*{-0.5cm} \\end{center} \\end{figure}  \\begin{table} \\caption{Model parameters.} \\label{par} \\vspace*{-0.3cm} \\begin{center} \\footnotesize \\begin{tabular}{c|cc|cc} \\tableline &\\multicolumn{2}{c|}{Model~A~(H$_\\circ\\!=\\!100 \\scriptstyle\\rm km\\,s^{\\!-\\!1}Mpc^{\\!-\\!1}$)} &\\multicolumn{2}{c}{Model~B~(H$_\\circ\\!=\\!69 \\scriptstyle\\rm km\\,s^{\\!-\\!1}Mpc^{\\!-\\!1}$)}\\\\ \\cline{2-5} Parameters& G1 & G2 & G1 & G2 \\\\ \\tableline s('') & 0.10 & 0.10       & 0.00 & 0.05 \\\\ b & 0.9072 & 0.2453     & 0.7797 & 0.3429 \\\\ e & 0.3269 & 0.6405        & 0.1570  & 0.3149 \\\\ $\\phi(^\\circ)$ & 163.45 & 154.93 & 172.62 & 160.86 \\\\ \\tableline $\\gamma,\\varphi_{\\gamma}(^\\circ)$ & 0.0876 & -10.92& 0.0473 & 12.47 \\\\ \\tableline \\end{tabular} \\end{center} \\vspace*{-0.4cm} \\caption{Comparison between observations and model predictions.} \\label{obs} \\vspace*{-0.3cm} \\begin{center} \\footnotesize \\begin{tabular}{cccc} \\tableline Properties & \\multicolumn{1}{c}{Observation} & \\multicolumn{1}{c}{Model A} & \\multicolumn{1}{c}{Model B} \\\\ \\tableline $\\vec x_A('')$ & (~0.0000,~0.0000) $\\pm$ (0.0023,0.0023) & (~0.0000,~0.0000) & (~0.0000,~0.0000)\\\\ $\\vec x_B('')$ & (-0.7380,-1.9612) $\\pm$ (0.0043,0.0046) & (-0.7382,-1.9613) & (-0.7365,-1.9518)\\\\ $\\vec x_C('')$ & (-0.7446,-0.4537) $\\pm$ (0.0045,0.0049) & (-0.7443,-0.4544) & (-0.7422,-0.4575)\\\\ $\\vec x_D('')$ & (~1.1284,-1.2565) $\\pm$ (0.0107,0.0124) & (~1.1271,-1.2582) & (~1.1269,-1.2207)\\\\ \\tableline $F_A / F_B$ & 2.042 $\\pm$ 0.124 & 1.917 & 1.901 \\\\ $F_C / F_B$ & 1.037 $\\pm$ 0.083 & 1.092 & 1.131 \\\\ $F_D / F_B$ & 0.350 $\\pm$ 0.055 & 0.428 & 0.504 \\\\ \\tableline $T_{AB} (d)$ & 26.0 $\\pm$ 5.0 & 28.4 & 27.6 \\\\ $T_{CB} (d)$ & 33.0 $\\pm$ 5.0 & 32.0 & 31.7 \\\\ $T_{DB} (d)$ & 73.0 $\\pm$ 5.0 & 68.4 & 71.3 \\\\ \\tableline $\\chi^2$ & 0.0 & 4.0 & 24.7 \\\\ \\tableline \\end{tabular} \\end{center} \\end{table}   \\vspace*{-1.0cm}"
    },
    "0002/astro-ph0002528_arXiv.txt": {
        "abstract": "Coalescing compact binaries are thought to be involved in a wide variety of astrophysical phenomena. In particular, they are important sources of gravitational radiation for both ground-based and space-based laser-interferometer detectors, and they may be sources of supernova explosions or gamma-ray bursts. Mergers of two white dwarfs may produce neutron stars with peculiar properties, including perhaps millisecond radio pulsars sometimes accompanied by planets (as observed in PSR 1257+12). According to a widely held belief, the coalescence of two neutron stars should produce a rapidly rotating black hole surrounded by an accretion disk or torus, but this is by no means certain. This review paper focuses on the final hydrodynamic coalescence and merger of double neutron stars and double white dwarfs, and addresses the question of the nature of the final merger products. ",
        "introduction": "The coalescence and merging of two compact stars into a single object is a very common end-point of close binary evolution. Dissipation mechanisms such as friction in common envelopes, tidal dissipation, or the emission of gravitational radiation, are always present and cause the orbits of close binary systems to decay. This review will concentrate on the coalescence of compact binaries containing either two neutron stars (hereafter NS) or two white dwarfs (WD).  \\subsection{Double Neutron Stars}  Many theoretical models of gamma-ray bursts (GRBs) rely on coalescing NS binaries to provide the energy of GRBs at cosmological distances (e.g., Eichler et al.\\ 1989; Narayan, Paczy\\'nski, \\& Piran 1992; M\\'esz\\'aros \\& Rees 1992; for recent reviews see M\\'esz\\'aros 1999 and Piran 1999). The close spatial association of some GRB afterglows with faint galaxies at high redshifts may not be inconsistent with a NS binary merger origin, in spite of the large recoil velocities acquired by NS binaries at birth (Bloom, Sigurdsson, \\& Pols 1999; but see also Bulik \\& Belczynski 2000). Currently the most popular models all assume that the coalescence of two NS leads to the formation of a rapidly rotating black hole (BH) surrounded by a torus of debris. Energy can then be extracted either from the rotation of the Kerr BH or from the material in the torus so that, with sufficient beaming, the gamma-ray fluxes observed from even the most distant GRBs can be explained (M\\'esz\\'aros, Rees, \\& Wijers 1999). However, it is important to understand the hydrodynamic processes taking place during the final coalescence before making assumptions about its outcome. In particular, as will be argued below (\\S2.2), it is not clear that the coalescence of two $1.4\\,M_\\odot$ NS forms an object that will collapse to a BH on a dynamical timescale, and it is not certain either that a significant amount of matter will be ejected during the merger to form an outer torus around the central object (Faber \\& Rasio 2000).  Coalescing NS binaries are also important sources of gravitational waves that may be directly detectable by the large laser interferometers currently under construction, such as LIGO (Abramovici et al.\\ 1992; see Barish \\& Weiss 1999 for a recent pedagogical introduction) and VIRGO (Bradaschia et al.\\ 1990). In addition to providing a major new confirmation of Einstein's theory of general relativity (GR), including the first direct proof of the existence of black holes (see, e.g., Flanagan \\& Hughes 1998; Lipunov, Postnov, \\& Prokhorov 1997), the detection of gravitational waves from coalescing binaries at cosmological distances could provide accurate independent measurements of the Hubble constant and mean density of the Universe (Schutz 1986; Chernoff \\& Finn 1993; Markovi\\'c 1993). Expected rates of NS binary coalescence in the Universe, as well as expected event rates in laser interferometers, have now been calculated by many groups. Although there is some disparity between various published results, the estimated rates are generally encouraging (see Kalogera 2000 for a recent review).  Many calculations of gravitational wave emission from coalescing binaries have focused on the waveforms emitted during the last few thousand orbits, as the frequency sweeps upward from $\\sim10\\,$Hz to $\\sim300\\,$Hz. The waveforms in this frequency range, where the sensitivity of ground-based interferometers is highest, can be calculated very accurately by performing high-order post-Newtonian (PN) expansions of the equations of motion for two {\\it point masses\\/} (see, e.g., Owen \\& Sathyaprakash 1999 and references therein). However, at the end of the inspiral, when the binary separation becomes comparable to the stellar radii (and the frequency is $\\go1\\,$kHz), hydrodynamics becomes important and the character of the waveforms must change. Special purpose narrow-band detectors that can sweep up frequency in real time will be used to try to catch the last $\\sim10$ cycles of the gravitational waves during the final coalescence (Meers 1988; Strain \\& Meers 1991). These ``dual recycling'' techniques are being tested right now on the German-British interferometer GEO 600 (Danzmann 1998). In this terminal phase of the coalescence, when the two stars merge together into a single object, the waveforms contain information not just about the effects of GR, but also about the interior structure of a NS and the nuclear equation of state (EOS) at high density. Extracting this information from observed waveforms, however, requires detailed theoretical knowledge about all relevant hydrodynamic processes. If the NS merger is followed by the formation of a BH, the corresponding gravitational radiation waveforms will also provide direct information on the dynamics of rotating core collapse and the BH ``ringdown'' (see, e.g., Flanagan \\& Hughes 1998).  \\subsection{Double White Dwarfs}  Coalescing WD binaries have long been discussed as possible progenitors of Type~Ia supernovae (Iben \\& Tutukov 1984; Webbink 1984; Paczy\\'nski 1985; see Branch et al.\\ 1995 for a recent review). To produce a supernova, the total mass of the system must be above the Chandrasekhar mass. Given evolutionary considerations, this requires two C-O or O-Ne-Mg WD. Yungelson et al.\\ (1994) showed that the expected merger rate for close WD pairs with total mass exceeding the Chandrasekhar mass is consistent with the rate of Type~Ia supernovae deduced from observations. Alternatively, a massive enough merger may collapse to form a rapidly rotating NS (Nomoto \\& Iben 1985; Colgate 1990). Chen \\& Leonard (1993) speculated that most millisecond pulsars in globular clusters might have formed in this way. In some cases planets may also form in the disk of material ejected during the coalescence and left in orbit around the central pulsar (Podsiadlowski, Pringle, \\& Rees 1991). Indeed the very first extrasolar planets were discovered in orbit around a millisecond pulsar, PSR B1257$+$12 (Wolszczan \\& Frail 1992). A merger of two magnetized WD might lead to the formation of a NS with extremely high magnetic field, and this scenario has been proposed as a source of GRBs (Usov 1992).  Close WD binaries are expected to be extremely abundant in our Galaxy, even though their direct detection remains very challenging (Han 1998; Saffer, Livio, \\& Yungelson 1999). Iben \\& Tutukov (1984, 1986) predicted that $\\sim20$\\% of all binary stars produce close WD pairs at the end of their stellar evolution. More recently, theoretical estimates of the double WD formation rate in the Galaxy have converged to a value $\\simeq0.1\\,{\\rm yr}^{-1}$, with an uncertainty that may be only a factor of two (Han 1998; Kalogera 2000). The most common systems should be those containing two low-mass helium WD. Their final coalescence can produce an object massive enough to start helium burning. Bailyn (1993 and references therein) and others have suggested that some ``extreme horizontal branch'' stars in globular clusters may be such helium-burning stars formed by the coalescence of two WD. Planets in orbit around a massive WD may also form following the binary coalescence (Livio, Pringle, \\& Saffer 1992).  Coalescing WD binaries are also important sources of low-frequency gravitational waves that should be easily detectable by future space-based laser interferometers. The currently planned LISA (Laser Interferometer Space Antenna; see Folkner 1998) should have an extremely high sensitivity (down to a characteristic strain $h\\sim10^{-23}$) to sources with frequencies in the range $\\sim10^{-4}\\,-\\,1\\,$Hz. Han (1998) estimated a WD merger rate $\\sim0.03\\,{\\rm yr}^{-1}$ in our own Galaxy. Individual coalescing systems and mergers may be detectable in the frequency range $\\sim10$--$100\\,$mHz. In addition, the total number ($\\sim10^4$) of close WD binaries in our Galaxy emitting at lower frequencies $\\sim0.1$--$10\\,$mHz (the emission lasting for $\\sim10^2$--$10^4\\,$yr before final merging) should provide a continuum background signal of amplitude $h\\sim10^{-20}$--$10^{-21}$ (Hils et al.\\ 1990). The detection of the final burst of gravitational waves emitted during an actual merger would provide a unique opportunity to observe in ``real time'' the hydrodynamic interaction between the two degenerate stars, possibly followed immediately by a supernova explosion, nuclear outburst, or some other type of electromagnetic signal. ",
        "conclusions": ""
    },
    "0002/astro-ph0002144_arXiv.txt": {
        "abstract": "Theoretical aspects of modeling of spectra of late M- and L-dwarfs are discussed. We show, that the processes of formation of spectra of M- and L-dwarfs are basically different. Instead of the case of M-dwarfs, atoms of Ti and V should be depleted into grains in the atmospheres of L-dwarfs. Overall shape of the L-dwarf spectra is governed by the K I + Na I resonance line wings of the huge strength. To fit lithium lines observed in spectra of the coolest dwarfs we used two additional suggestions: a) there are some {\\em extra} depletions of molecular species absorbed in the optical spectra of L-dwarfs; b) there may be (a few?) additional (``dusty''?) opacity sources in their atmospheres. Problems of lithium line formation and the ``natural'' limitation of their use for the ``lithium test'' for the case of L-dwarfs are considered. ",
        "introduction": "A few definitions of the ``unconventional'' spectral classification of low mass stars and substellar objects are used in this paper:   \\\\  {\\bf M-dwarfs~} are objects (stars + young brown dwarfs) with 4000 $> T_{\\rm eff} >$ 2200 K. Their spectra are governed by molecular bands of TiO and VO (in the optical part of the spectrum) , H$_2$O (in the red).  \\\\  {\\bf L-dwarfs~} are objects (brown dwarfs + stars) with 1000 $<$ $T_{\\rm eff}$ $<$  2200 K (cf. Martin et al. 1997). Optical spectra of L-dwarfs are governed by the K I + Na I lines and molecular bands of CrH + FeH (in the optical spectrum), H$_2$O (in the red).  \\\\  {\\bf T-dwarfs~} are super-giant, planet-like objects (big Jupiters?) with T$_{\\rm eff} <$ 1000 K. Their spectra are formed by the \"dusty opacities\" and methane + H$_2$O + ...? absorption (Strauss et al. 1999, Burgasser et al. 1999)\\\\  For the time being most of the known  brown dwarfs are actually recognized by the detection of  the  Li\\,{\\sc i}  resonance doublet  in  their  spectra   (Rebolo et al. 1996, Mart\\'\\i n et al. 1997a,  Kirkpatrick  et  al. 1999). Indeed, temperatures in the interiors of brown dwarfs are not high enough to burn lithium (Rebolo et al. 1992). ",
        "conclusions": "Finally, we have arrived to the following conclusions:  \\begin{itemize}  \\item Processes of formation of Li lines in L- and M-spectra differ significantly:  -- for M-dwarfs the main problem to be solved is blending of lithium lines by molecular lines,  -- in the case of L-dwarfs we deal with a menagerie of different processes: depletion of lithium atoms into molecules and grains, \"dusty effects\", meteorological phenomena, stratification effects, etc...  \\item We can fit the optical spectra of L-dwarfs in the frame  of our simple model.  \\item  Using our model we may perform a numerical analysis of the L-dwarf spectra (at least in the sense of the Li abundance determination).  \\item The basic algorithm of the ``lithium test'' may be used even for the assessment of the coolest L-dwarfs.  \\end{itemize}"
    },
    "0002/astro-ph0002372_arXiv.txt": {
        "abstract": "The luminosity function of MS2255.7+2039 at $z=0.288$ is determined down to a total magnitude of $R \\sim 24$, corresponding to $M_{\\mbox{\\tiny{\\rm R}}} \\sim -17 \\, ({\\rm H}_{\\mbox{\\tiny{0}}} = 50 \\ {\\rm km \\ s}^{-1} \\ {\\rm Mpc}^{-1})$. The data are corrected for incompleteness and crowding using detailed simulations. We find that the luminosity function is steeper than a standard Schechter function at faint magnitudes, and shows an excess of galaxies below $ M_{\\mbox{\\tiny{\\rm R}}} \\sim -19$. After corrections for light loss and crowding, the data can be described by a sum of two Schechter functions, one with $M_{\\mbox{\\tiny{\\rm R}}}^{*} = -22.8$ and $\\alpha = -1.0$, and one steeper with $M_{\\mbox{\\tiny{\\rm R}}}^{*} = -18.9$ and $\\alpha = -1.5$, representing the dwarf population. A straight-line fit to the faint part yields a slope similar to the Schechter $\\alpha = -1.5$ of the dwarf population. The luminosity function of MS2255.7+2039 is compared to other clusters at lower redshifts, and does not show any significant difference. The redshift range for clusters in which increased number of dwarf galaxies have been found is therefore extended to higher redshifts. ",
        "introduction": "\\label{sectintrod}  The faint end of the luminosity function (LF) is of great interest both in connection to the excess found in number counts of faint galaxies and for theories of cluster formation. The latter is demonstrated especially in a series of papers by  Kaufmann et al. (1993), Kaufmann (1995a, b), Heyl et al. (1995) and Baugh et al. (1996), which show how the LF is related to the hierarchical clustering in e.g. the Cold Dark Matter (CDM) model. A steep faint end of the LF is unique not only to the CDM scenario, but is a generic prediction of hierarchical models of cluster formation (White \\& Frenk 1991). The difference of the LF in the field and in clusters of galaxies also  gives important information about environmental effects related to the galaxy formation process. In particular, ram pressure stripping and effects of interaction are likely to be more important for clusters than in the field (e.g., Moore et al. 1996). There is in this respect a clear connection to the Butcher-Oemler effect (Butcher \\& Oemler 1984), seen for clusters, but not apparent in the field. The fraction of blue galaxies in clusters is larger at higher redshifts, which is often interpreted as an evolutionary effect. The evolution of the cluster LF is therefore of obvious interest.  The field LF has been investigated locally by several groups. Some of these report local LFs with flat faint-end slopes of $-1.1 \\la \\alpha \\la -1.0$ (Loveday et al. 1992; Ellis et al. 1996), while others have found a significantly steeper LF with $\\alpha \\la -1.5$ (Marzke et al. 1994; Lilly et al. 1995). At higher redshifts especially the CFRS survey (Lilly et al. 1995) and the Autofib survey (Ellis et al. 1996) give information about the LF in the field up to $z \\sim 1$. Lilly et al. find that, while strong evolution is seen for the blue sample at $z \\gsim 0.5$, the red LF is  changing little. At $z \\sim 0.5$  the LF has brightened by $\\sim 1$ magnitude in B. At higher redshifts the bright end stays constant, while the faint continues to increase, leading to a steepening of the faint end of the LF, from $\\alpha = -1.1$ locally to $\\alpha = -1.5$ at $z \\simeq 0.5$, in broad agreement with the Autofib survey. A steepening of the LF can explain the excess counts found in deep surveys, as discussed by Gronwall \\& Koo (1995).  While the field LF has been studied in fair detail, there have been only a few CCD investigations of the LF of clusters of galaxies until recently. Relatively nearby clusters, like Virgo, Fornax and Coma, have been studied by a number of authors (e.g., Bernstein et al. 1995, hereafter BNTUW; Lobo et al. 1997; Biviano et al. 1995; De Propris et al. 1995; Trentham 1998a). While some earlier investigations yield faint-end slopes of $\\alpha \\simeq -1.3$ (e.g., Ferguson \\& Sandage 1988), more recent investigations, taking low surface brightness galaxies into account, point to steeper slopes, $\\alpha \\la -1.5$ (e.g., Bothun et al. 1991).   At higher redshifts the information about the cluster LF is scarce. The first detailed study was that by Driver et al. (1994b, henceforth DPDMD), who studied the R-band LF to $R = 24$ for the cluster Abell 963 at $z = 0.206$. While the high luminosity end can be well fitted with a Schechter function, they found an increase of faint galaxies between $-19 \\lsim M_{\\mbox{\\tiny{\\rm R}}} \\lsim -16.5$, with a slope of $\\alpha \\simeq -1.8$. Further investigations at $0.1 \\la z \\la 0.2$ have strengthened the case for steep slopes, $\\alpha \\la -1.7$, at $-19 \\la M_{\\mbox{\\tiny{\\rm R}}}$ (Smith et al. 1997; Wilson et al. 1997, hereafter WSEC), although examples of flatter LFs also exist (Trentham 1998b).  It is obvious that these results need confirmation for more clusters, especially at higher redshifts. Because of the strong evolution of the Butcher-Oemler effect, as well as from numerical simulations of cluster evolution (e.g., Kauffmann 1995a,b), one expects substantial evolution even from $z = 0.2$ to $z = 0.4$. The detection of the faint end of the LF then becomes more difficult, both because the galaxies are fainter, and because of increasing contamination from the background. As a first step we here report observations of  the cluster MS2255.7+2039 at $z = 0.288$. We will in this paper use ${\\rm H}_{\\mbox{\\tiny{0}}} = 50 \\kms~Mpc^{-1}$ and $\\Omega_{\\mbox{\\tiny{\\rm M}}} = 1$. ",
        "conclusions": "We have observed the galaxy cluster MS2255.7+2039 ($z = 0.288$) and a background field at similar galactic latitude with the aim of determining the cluster LF. The isophotal magnitudes have been corrected for light loss according to results obtained from simulations. We have also compensated for obscuration due to bright, apparently large, objects in the images. The resulting cluster LF has a fairly steep faint-end slope ($\\alpha \\simeq -1.5$) faintward of the break in the profile around $M_{\\mbox{\\tiny{\\rm R}}} = -19$. This slope is more shallow than some LFs found in clusters both locally and at $z \\simeq 0.2$, but similar to the slope of the composite LF derived by Trentham (1998c). Without focusing too strongly on the precise value of the slope, we conclude that MS2255.7+2039 exhibits a steepening LF at faint magnitudes.  The evidence for steep faint-end slopes of cluster LFs is accumulating. There are now a number of fairly deep CCD investigations of nearby, as well as a few medium distant ($z \\la 0.3$) clusters, that all point to rising LFs at faint magnitudes. It therefore seems clear that a flat LF (i.e., $\\alpha = -1$) can be ruled out even at intermediate magnitudes ($ -20 \\la M_{\\mbox{\\tiny{\\rm R}}} \\la -17$). However, several questions remain unanswered. The uncertainties in the measured slopes are probably considerable, since different correction methods seem to yield deviating results, which probably explains the discrepancies between the LFs found for the same clusters as determined by different investigators (see section~\\ref{comparison}).  Because of these uncertainties, it is too early to discuss any variation of the faint-end slope with $z$. The accuracy of the present study only allows us to claim that the cluster LF is non-flat at faint magnitudes ($-19 \\la M_{\\mbox{\\tiny{\\rm R}}}$). The exact values of the slope and the magnitude where the steepening sets in are uncertain, and any trend with $z$ that may be present is dominated by these uncertainties. In addition, environmental differences, like richness or density, between clusters at the same $z$ could affect individual LFs, making distinctions in $z$ even more difficult to isolate. The uncertainties in the background subtraction is also a source of error, although the simulations by Driver et al. (1998) show that the faint-end slope of the LF can be reliably determined out to $z \\simeq 0.3$ with seeing and depth similar to those of the present data. Nevertheless, the errors due to the statistical background subtraction can probably be substantially reduced by using photometric redshifts. Work based on this approach is in progress.  There are in the context of cluster LFs several important questions to answer in the future: Is there a universal LF for galaxy clusters at low redshifts, or is the steepness of the dwarf population different between clusters? Is there a colour dependence of the steepness of the dwarf population, as may be indicated in the study by WSEC? The K-band observations of five clusters by Trentham \\& Mobasher (1998) are especially interesting in this context. These data were, however, not deep enough to draw any conclusions about the faint end of the luminosity function. From the clear signs of evolution between $z = 0.5$ and the present epoch for the field (e.g., Lilly et al. 1995), one would expect a corresponding evolution in clusters. From the galaxy harassment scenario for the Butcher-Oemler effect (Moore et al. 1996) one may expect a larger fraction of low luminosity galaxies in the past, and therefore a steeper LF. We hope to address some of these issues in the future."
    },
    "0002/astro-ph0002067_arXiv.txt": {
        "abstract": "{Reconnection is the process by which magnetic fields in a conducting fluid change their topology. This process is essential for understanding a wide variety of astrophysical processes, including stellar and galactic dynamos and astrophysical turbulence. To account for solar flares, solar cycles and the structure of the galactic magnetic field reconnection must be fast, propagating with a speed close to the Alfv\\'en speed.  Earlier attempts to deal with magnetic reconnection assumed that magnetized fluids are laminar and as a result obtained slow reconnection rates.  We show that the presence of a random magnetic field component substantially enhances the reconnection rate and enables {\\it fast reconnection}, i.e. reconnection that does not depend on fluid resistivity. The enhancement of the reconnection rate is achieved via a combination of two effects.  First of all, only small segments of magnetic field lines are subject to direct Ohmic annihilation. Thus the fraction of magnetic energy that goes directly into fluid heating goes to zero as fluid resistivity vanishes. However, the most important enhancement comes from the fact that unlike the laminar fluid case where reconnection is constrained to proceed line by line, the presence of turbulence enables many magnetic field lines to enter the reconnection zone simultaneously. A significant fraction of magnetic energy goes into MHD turbulence and this enhances reconnection rates through an increase in the field stochasticity. In this way magnetic reconnection becomes fast when field stochasticity is accounted for. As a consequence solar and galactic dynamos are also fast, i.e. do not depend on fluid resistivity.} ",
        "introduction": " ",
        "conclusions": "It is not possible to understand the dynamics of magnetized astrophysical plasmas without understanding how magnetic fields reconnect. Here we have compared traditional approaches to the problem of magnetic reconnection and a new approach that includes the presence of turbulence in the magnetized plasma.  One of the more striking aspects of our result is that the global reconnection speed is relatively insensitive to the actual physics of reconnection.  Equation (\\ref{main}) only depends on the nature of the turbulent cascade.  Although this conclusion was reached by invoking a particular model for the strong turbulent cascade, we showed in LV99 that any sensible model gives qualitatively similar results. One may say that the conclusion that reconnection is fast, even when the local reconnection speed is slow, represents a triumph of global geometry over the slow pace of ohmic diffusion.  In the end, reconnection can be fast because if we consider any particular flux element inside the contact volume, assumed to be of order $L_x^3$, the fraction of the flux element that actually undergoes microscopic reconnection vanishes as the resistivity goes to zero.  The new model of fast turbulent reconnection changes our understanding of many astrophysical processes.  Firstly, it explains why dynamos do not suppress themselves through the excessive generation of magnetic noise, as some authors suggest (Kulsrud \\& Anderson 1992). The model also explains why reconnection may be sometimes fast and sometimes slow, as solar activity demonstrates.  ADAFs and the acceleration of cosmic rays at reconnection sites are other examples of processes where a new model of reconnection should be applied.  Our results on turbulent reconnection assume that the turbulent cascade is limited by plasma resistivity. If gas is partially ionized collisions with neutrals may play an important role in damping turbulence.  A study in Lazarian \\& Vishniac (2000) shows that for gas with low levels of ionization turbulent reconnection may be impeded as magnetic field wandering is suppressed on small scales. However, the level of suppression depends on the details of the energy injection into the turbulent cascade (see a discussion in Lazarian \\& Pogosyan 2000), which are far from being clear. Moreover, for very low ionization levels there will be an enhancement of the reconnection process as neutrals diffuse perpendicular to magnetic field lines. Thus, reconnection may still be an important process in the evolution of molecular clouds and in star formation."
    },
    "0002/astro-ph0002298_arXiv.txt": {
        "abstract": "1ES~1101--232 is a bright BL Lac of the High frequency peak class. We present here the results of two \\sax\\ observations in which the source has shown a variation of about 30\\% in flux with a corresponding spectral variability. We interpret the overall spectral energy distribution in terms of an homogeneous SSC model and, by using also the TeV upper limit from a short Mark 6 pointing, derive constraints on the physical parameters of the source, in particular on the magnetic field strength. The overall Spectral Energy Distribution makes 1ES~1101--232 a very promising candidate for TeV detection. ",
        "introduction": "\\noindent  BL Lac objects form a minority class of active nuclei (see e.g. Urry \\& Padovani, 1995), but nevertheless their high luminosities and extreme variability in all bands make them an interesting subclass to study intrinsic properties of nuclear emission. Furthermore, it is matter of discussion the position of BL Lacs with respect to other active nuclei in the framework of Unification Models. The current picture claims that BL Lacs are the fraction of Fanaroff-Riley I galaxies that point their jet towards us, but many details still need to be worked out. Often, BL Lacs are studied together with other classes of flat radio spectrum sources, to form the class of blazars.  Recently a sequence has been proposed for the subclass of gamma ray bright blazars, and possibly valid for all blazars, based on their bolometric luminosity (Ghisellini et al. 1998, Fossati et al. 1998). The overall spectral energy distribution has two peaks (in $\\nu F_\\nu$ representation), the one at lower energies due to synchrotron radiation and the higher energy one to Inverse Compton scattering. It is proposed that the value of the peak frequency is the result of the balance between radiative cooling and acceleration of the corresponding electrons (see e.g. Ghisellini, 1999): in less luminous objects the radiative cooling is less efficient, allowing the accelerated electrons to reach higher energies. As a consequence, both the synchrotron and the inverse Compton peaks shift to higher frequencies as the bolometric intrinsic luminosity decreases. In this view, the class of BL Lacertae objects known as HBL (High frequency peak BL Lacs) has a smaller bolometric luminosity and a higher synchrotron peak frequency than the class of LBL (Low frequency peak BL Lacs). HBL are mostly found in X--ray selected samples, since they are expected to produce most of their synchrotron emission in the X--ray band. If the emission peak (as in the LBL objects) is at frequencies smaller than the observed X--ray band, the X--ray spectrum is steep (i.e. the steepening part above the peak), while HBL with a flat X--ray spectrum should have their peak in the observed X--ray band.  As part of a program aimed at a spectral survey of soft X--ray selected BL Lacs with \\sax\\ (Wolter et al. 1998), we have studied also 1ES~1101--232 (z=0.186), a bright BL Lac selected from the Slew Survey (Perlman et al. 1996) that shows an extreme behavior with a very flat X--ray spectrum. The object was detected, besides in the 0.1-10 keV energy band, also up to $\\sim$ 100 keV in only $\\sim 6000$ sec. The best fit of the source in the 0.1-100 keV band is given by a broken power law, with Galactic low energy absorption, that has a break energy E$_0$=1.36 (1.11--1.65) keV. The spectral slope (energy index) derived from the PDS is consistent with the one derived from the MECS above the break energy E$_0$: $\\alpha_x$ = 1.03 (0.99--1.08) (Wolter et al. 1998).  We have constructed also the Spectral Energy Distribution (SED) for this object, using flux measurements collected from the literature, from radio to X--rays. We have fitted a cubic polynomial to the distribution in order to find the peak of the SED, that indeed falls in the \\sax\\ band (log $\\nu_{peak} \\sim$ 17.48, corresponding to $\\sim 1.3$ keV). For this object, therefore, the break energy derived from the spectral fit in the X--ray band is consistent, albeit within its large indetermination, with the position of the synchrotron peak as derived from the overall distribution (SED).  The SED of 1ES~1101--232 is similar to that of the flaring state of Mkn 501 (Pian et al. 1998) and 1ES~2344+514 (Catanese et al. 1998; Giommi et al. 2000), and therefore we could expect a strong TeV emission. A quasi--simultaneous X--ray and TeV observation has been therefore scheduled, to confirm the X--ray spectrum of the source, and its overall shape, to detect possible variations in flux, that could constrain the physical parameters of the source, and to monitor the TeV emission to search for possible detection.  \\begin{table*} \\caption{Fit results for 1ES~1101--232 LECS+MECS data OBS. 19/06/98 {\\bf Low state}} \\begin{tabular}[h]{| l l c c c r r r|} \\hline Model  &   $N_{\\rm H}^a$ & $\\alpha_1$ & $\\alpha_2$   & $E_0$     & $F^{\\rm b}$ &  $\\chi^2(dof)$ & Prob. \\\\ &           &               &                 & keV       &             &  & \\\\ (1)  & 5.76        &    --         & 1.19(1.16-1.22) & --        & 2.60     & 267.1(195) & $<$ 0.5\\%\\\\ (2)  & 8.3(7.1-10.4) &  --         & 1.25(1.22-1.29) & --        & 2.57     & 240.5(195) & 1.5\\%\\\\ (3)  & 5.76        &0.80(0.58-0.96)& 1.29(1.25-1.33) & 1.34(1.08-1.82) &2.55 & 216.2(194) & 13\\% \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item (1) Single p.l., $N_{\\rm H}$=$N_{\\rm H}^{\\rm Gal}$, LECS/MECS ratio free; (2) Single p.l., $N_{\\rm H}$ free, LECS/MECS ratio=0.7 ; (3) Broken p.l., $N_{\\rm H}=N_{\\rm H}^{\\rm Gal}$, LECS/MECS ratio = 0.7. Errors quoted are 90\\% confidence intervals. \\\\ $^{\\rm a}$ Column density in $\\e20$. \\\\ $^{\\rm b}$ Unabsorbed [2--10 keV] flux in $10^{-11} \\ecs$. \\\\ \\end{list} \\label{broken_n} \\end{table*} \\begin{table*} \\caption{Fit results for 1ES~1101--232 LECS+MECS data OBS. 04/01/97 {\\bf High state} } \\begin{tabular}[h]{| l l c c c r r r|} \\hline Model  &   $N_{\\rm H}^{\\rm a}$ & $\\alpha_1$ & $\\alpha_2$ & $E_0$  & $F^{\\rm b}$ &  $\\chi^2(dof)$ & Prob. \\\\ &           &               &                 & keV        &             &  & \\\\ (1)  & 5.76          &  --          & 0.97(0.95-1.00) & --        & 3.81       & 224.0 (182) & $<$ 2\\%\\\\ (2)  & 8.9(7.2-12.6) &  --          & 1.03(0.99-1.08) & --        & 3.79       & 205.1 (182) &  11 \\%\\\\ (3)  & 5.76      &0.59(0.35-0.74)   & 1.05(1.01-1.08) & 1.36(1.11-1.65) &3.76  & 191.6 (181) & 28 \\%\\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item (1) Single p.l., $N_{\\rm H}$=$N_{\\rm H}^{\\rm Gal}$, LECS/MECS ratio free; (2) Single p.l., $N_{\\rm H}$ free, LECS/MECS ratio=0.7; (3) Broken p.l., $N_{\\rm H}$=$N_{\\rm H}^{\\rm Gal}$, LECS/MECS ratio = 0.7. Errors quoted are 90\\% confidence intervals. \\\\ $^{\\rm a}$ Column density in $\\e20$. \\\\ $^{\\rm b}$ Unabsorbed [2-10 keV] flux in $10^{-11} \\ecs$. \\\\ \\end{list} \\label{broken_o} \\end{table*}   The \\sax\\ data are presented here. The TeV observation, conducted in non--optimal weather condition, did not yield a detection, but we will use the upper limit (Chadwick et al. 1999a) to derive useful information on the physical mechanisms at work in this source. The plan of the paper is as follows: in Section 2 we describe the observational data obtained with \\sax\\ in the two epochs, in Section 3 we summarize the TeV predictions and observations, that help constraining the parameters of the source, by using the SED and theoretical models of emission as explained in Section 4. Section 5 presents our results and conclusions.   Throughout the paper a Hubble constant H$_0$=50 $km\\, s^{-1}\\, Mpc^{-1}$ and a Friedman universe with a deceleration parameter $q_0$=0 are assumed. ",
        "conclusions": "The X--ray spectrum of 1ES~1101--232 as measured by \\sax\\ is fitted only by a broken power law (a single power law or an absorbed power law are not statistically acceptable) with a break at 1.3 - 1.9 keV. From the first to the second observation, the spectrum varied at high energies, becoming softer (steeper). The flux decrease, by about 32\\%, has occurred in the 2--10 keV band. The PDS observation are not of statistical significance sufficient to put a real constraint on the spectrum.  Even if the variation in the X--ray band is not dramatic, we can clearly distinguish between the two states, that are modeled by different parameters of the synchrotron component.  By using the TeV upper limit and the two \\sax\\ observations we model also the higher energy portion of the spectrum as a self--Compton component, by using the model described e.g. in Tavecchio et al. (1998) that assumes a simple homogeneous SSC model in the KN regime, in which the relativistic electrons have a broken power law energy distribution.  The two X-ray states of the source are described by varying this distribution, assuming that the other relevant parameters ($R$ and $\\delta$) are nearly constant.   We can compare these results with what found for the few TeV detected sources. The choice of Doppler factor of 10 and 20 made here is in the interval of the values of $\\delta$ found by other authors for Mkn 421, PKS2155--304 and Mkn501, that range from $\\delta \\sim 5$ (e.g. Takahashi, 1999; Mkn 421, Catanese et al. 1998: Mkn 501) to $\\delta \\sim 30$ (e.g. Bednarek \\& Protheroe 1999: Mkn501; Kataoka et al 2000: PKS2155--304; Bednarek \\& Protheroe 1997: Mkn 421). At the same time, values of $B$ in excess of 0.03 G and up to 1 G (Chiappetti et al. 1999) are found for the same sources.  Of the three above mentioned objects, the most similar to 1ES~1101--232 is Mkn 501, for which different measures have been produced: e.g. $\\delta \\geq 15$ \\& $B=0.8 $G (Pian et al. 1998); $\\delta = 15$ \\& $B=0.2 $G (Kataoka et al. 1999); $\\delta = 30$ \\& $B=0.7 $G (Bednarek \\& Protheroe 1999), in good agreement with the values of Table~\\ref{params}.   Even if based, besides the accurate X-ray spectral determination up to $\\sim 50$ keV, only on a TeV upper limit we can infer that the physical conditions in 1ES~1101--232 are similar to the brightest TeV sources, making it a very promising candidate for TeV observations, and a testbed for the SSC model. Furthermore, since the redshift of 1ES~1101--232 is intermediate, a detection of this source in the VHE range would pose constraints on the density of the IR background photons that is still at the moment very uncertain."
    },
    "0002/astro-ph0002251_arXiv.txt": {
        "abstract": "A new method for the inclusion of ionizing radiation from uniform radiation fields into 3D Smoothed Particle Hydrodynamics ({\\sc sphi}) simulations is presented. We calculate the optical depth for the Lyman continuum radiation from the source towards the {\\sc sphi} particles by ray-tracing integration. The time-dependent ionization rate equation is then solved locally for the particles within the ionizing radiation field. Using test calculations, we explore the numerical behaviour of the code with respect to the implementation of the time-dependent ionization rate equation. We also test the coupling of the heating caused by the ionization to the hydrodynamical part of the {\\sc sphi} code. ",
        "introduction": "Smoothed Particle Hydrodynamics ({\\sc sph}) has become a numerical method widely used for addressing problems related to fluid flows in astrophysics. Due to its Lagrangian nature it is especially well suited for applications involving variations by many orders of magnitude in density. Examples for this type of applications are simulations of the collapse of molecular clouds and the formation of a stellar cluster, as performed by Klessen, Burkert \\& Bate~\\shortcite{klebuba}. A comparison between grid based methods and {\\sc sph} was performed by Burkert, Bate \\& Bodenheimer~\\shortcite{bbb} and Bate \\& Burkert~\\shortcite{bateburk}. They applied both methods to the numerically demanding problem of gravitational collapse and fragmentation of a slightly perturbed rotating cloud core with an $r^{-2}$ density profile. Both methods yielded the same qualitative results. Bate~\\shortcite{bate2} performed the first calculation which followed the collapse of a molecular cloud core in 3 dimensions down to a protostellar object in hydrodynamical equilibrium, thus spanning 17 (!) orders of magnitude in density. Other applications include accretion processes in massive circumbinary disks \\cite{boba1,babo1}, the collapse of cloud cores induced by shock waves \\cite{vanhala} or colliding clumps \\cite{bhattal}, the precession of accretion disks in binary systems \\cite{larwood}, the dynamical behaviour of massive protostellar disks \\cite{nelson} or the formation of large scale structure and galaxies in the early universe \\cite{steinmetz}.  A variety of physical processes are at work in the interstellar medium, like magnetic fields, radiation or thermal conductivity, necessitating their inclusion into numerical codes. This has already been achieved to a large extent in grid based methods like the magneto-hydrodynamics codes {\\sc zeus} \\cite{stono} or {\\sc nirvana} \\cite{ziegler}, or codes including effects of IR and UV radiation \\cite{yorka,sonnha,richling}.  In contrast, the addition of physical processes to {\\sc sph} codes is just at its beginnings. Extensions achieved so far are sophisticated equations of state (e.g. Vanhala et al.~1998) and self-gravity. Some efforts were made to make {\\sc sph} faster and more accurate. The introduction of {\\sc tree} algorithms \\cite{barnes,press,benz} and the use of GRAvity PipE ({\\sc grape}), a hardware device for fast computation of the gravitational N-body forces \\cite{umemura,steinmetz}, helped reducing the numerical effort for the gravitational force calculation and the determination of the nearest neighbours for each particle. Inutsuka~\\shortcite{inutsuka} presented a Godunov-like solver for the Eulerian equations in {\\sc sph} thus enhancing the numerical treatment of shocks. The introduction of gravitational periodic boundaries~\\cite{hernqu,klessen} allows the treatment of fragmentation and turbulence in molecular clouds without global collapse. The timestep problem which arises during isothermal collapse calculations at high densities is circumvented by the formation of sink particles, which substitute the innermost parts of the collapsing clump by one particle and accumulate the infalling mass and momenta \\cite{babopri}.  The strength of {\\sc sph} lies in its Lagrangian nature, which makes it especially attractive for problems involving gravitational collapse and star formation. Applications like e.g. by Klessen et al.~\\shortcite{klessen}, which deal with the collapse and fragmentation of molecular clouds, neglect the feedback processes of newly born stars which act on their parental cloud through stellar winds, outflows and ionization. This simplification may be justified as long as the simulations deal with collapse on timescales smaller than $\\approx$ 1 Myr, on which single and binary stars or T Tau-like clusters are formed~\\cite{efremov}. The case is different for larger timescales, on which OB subgroups and associations are formed. Neglecting feedback in these cases can lead  to unphysical results, like a star formation efficiency of 100 per cent, since in the purely isothermal case all material will sooner or later be accreted onto the evolving protostellar cores. This is in strong contradiction to observations, which estimate a global star formation efficiency for ordinary molecular clouds of order 10 per cent \\cite{pppII}. Another possible effect of feedback is the induction of star formation due to the compression of cloud material by shock waves and ionization fronts.  In this paper we discuss the implementation of the effects of ionizing UV radiation by massive stars into {\\sc sph} calculations as a first step in order to perform collapse calculations on scales where OB-stars are formed in a more realistic way. This will in future applications allow us to assess questions like: How does the process of ionization by massive stars change the stellar initial mass function? What are the implications for the star formation efficiency? Can star formation be induced by ionization, and if yes, what are the time scales and the parameter space, for which induced star formation can be expected? These questions will be discussed in subsequent papers. ",
        "conclusions": "The method presented in this paper allows the treatment of the dynamical effects of ionizing radiation in {\\sc sph} calculations. Thus the study of astrophysical problems arising from ionization, like the impact of ionizing radiation from newly born stars onto the evolution of their parental molecular clouds or the more consistent treatment of heating by OB associations in galaxy dynamics calculations are now feasible for the first time with {\\sc sphi} in 3 dimensions. We demonstrate that the code is able to treat time-dependent ionization, the related heating effects and hydrodynamics correctly. Our first applications, detailed calculations of photoionization induced collapse in molecular clouds and results obtained from them, will be presented in a subsequent paper.  To allow the correct treatment of recombination zones, one has to include the effects of time dependent heating and cooling processes by ionization and recombination, emission of forbidden lines and thermal radiation from dust. Another important aspect which was neglected here is the effect of the diffuse Lyman continuum recombination field. It can lead to the penetration of regions shielded from the direct ionizing radiation by the ionization front, which is e.g. seen in calculations of photoevaporating protostellar disks \\cite{yowe,richling}. An implementation of these processes into our {\\sc sphi} code is planned in the future."
    },
    "0002/astro-ph0002121_arXiv.txt": {
        "abstract": "We present results from a study of 2--8~keV X-ray sources detected by the Advanced CCD Imaging Spectrometer (ACIS) instrument on the \\chandra\\ X-ray Observatory in the field of the $z=0.516$ cluster \\crsscluster. In our 63.5~arcmin$^2$ search area, we detect 10 sources with 2--8~keV fluxes down to $\\approx 4\\times 10^{-15}$~erg~cm$^{-2}$~s$^{-1}$; our lowest flux sources are $\\approx 10$ times fainter than those previously available for study in this band. Our derived source density is about an order of magnitude larger than previous source counts above 2~keV, although this density may be enhanced somewhat due to the presence of the cluster. We detail our methods for source detection and characterization, and we show that the resulting source list and parameters are robust. We have used the Marcario Low Resolution Spectrograph on the Hobby-Eberly Telescope to obtain optical spectra for several of our sources; combining these spectra with archival data we find that the sources appear to be active galaxies, often with narrow permitted lines, red optical continua or hard X-ray spectra. Four of the X-ray sources are undetected to $R=21.7$; if they reside in $L^\\star$ galaxies they must have $z>$~0.55--0.75 and hard X-ray luminosities of $L_{2-8}\\simgt 4\\times 10^{42}$~erg~s$^{-1}$. We detect all but one of our 2--8~keV sources in the 0.2--2~keV band as well. This result extends to significantly lower fluxes the constraints on any large, completely new population of X-ray sources that appears above 2--3~keV. ",
        "introduction": "About 70\\% of the extragalactic X-ray background in the soft 0.5--2~keV band has been resolved into discrete sources by pencil-beam surveys with the \\rosat\\ satellite (e.g., \\cite{has98}). The 0.5--2~keV source counts have reached a surface density of $\\approx 1000$~deg$^{-2}$ at a discrete source detection limit of $\\approx 10^{-15}$~erg~cm$^{-2}$~s$^{-1}$, and simple extrapolations suggest that all of the 0.5--2~keV extragalactic X-ray background will be resolved by a flux limit of $\\sim 2\\times 10^{-16}$~erg~cm$^{-2}$~s$^{-1}$ at which the surface density will be $\\sim 3000$~deg$^{-2}$. Optical identification programs (e.g., \\cite{schmidt98}) have established that type~1 Active Galactic Nuclei (AGN), such as Seyfert~1 galaxies and Quasi-Stellar Objects (QSOs), are the dominant contributors above a 0.5--2~keV flux of $\\approx 5\\times 10^{-15}$~erg~cm$^{-2}$~s$^{-1}$. A non-negligible number (about 16\\%) of type~2 AGN are seen as well.  Largely due to instrumental limitations, the nature of the sources that produce the $>2$~keV X-ray background is much less certain at present. It is important to solve this mystery since most of the energy density in the X-ray background is located above the \\rosat\\ band. The best current constraints on the sources of the 2--10~keV X-ray background have come from the \\asca\\ and \\sax\\ satellites. Long observations with these satellites have reached discrete source detection limits of $\\approx$~(3--5)$\\times 10^{-14}$~erg~cm$^{-2}$~s$^{-1}$ and have resolved $\\approx 30$\\% of the 2--10~keV background into discrete sources (e.g., Ogasaka et~al. 1998; Giommi, Fiore \\& Perri 1999). The integrated number of sources, $N$, is consistent with the law $N(>S)\\propto S^{-3/2}$ expected for a uniform distribution of sources in Euclidean space. The deepest 2--10~keV source counts to date have resolved $\\approx 60$ sources deg$^{-2}$. The faintest 2--10~keV sources appear to have flatter spectra (with energy indices of $\\alpha\\approx 0.5\\pm 0.2$) than those of typical unabsorbed AGN (e.g., Ueda et~al. 1998), suggesting that a population of sources with spectra similar to that of the integrated X-ray background dominates above 2~keV. The population is thought to be at least partially composed of obscured AGN, and some of these hard sources have indeed been associated with such objects (e.g., Fiore et~al. 1999; Akiyama et~al. 2000). An important result, however, is that the majority of the hard sources found thus far appear to have counterparts in the soft X-ray band (Giommi, Fiore \\& Perri 1998). Obscured AGN might still create soft X-ray emission via electron-scattered X-rays or due to non-nuclear X-ray emission (e.g. starburst activity).  The arcsecond imaging quality and high-energy sensitivity of the \\chandra\\ X-ray Observatory (\\cite{wod96}) promises to revolutionize our understanding of the X-ray background above 2~keV. The source confusion and misidentification problems that have dogged earlier hard X-ray surveys will be eliminated. In this paper, we use data from \\chandra, the 8-m class Hobby-Eberly Telescope (HET), and public archives to study a small but well-defined sample of faint X-ray sources in the 2--8~keV band. Our sample contains several of the faintest $>2$~keV sources yet detected and identified. We address (1) the number counts at faint hard X-ray fluxes, (2) the nature of the faint hard X-ray sources, and (3) the issue of whether most faint hard X-ray sources have soft X-ray counterparts. Our X-ray sources lie in the PG~0027+260 (an eclipsing cataclysmic variable) field of the Cambridge-Cambridge \\rosat\\ Serendipity Survey (e.g., Boyle, Wilkes \\& Elvis 1997). This field contains \\crsscluster, a $z=0.516$ cluster of galaxies, and it was observed by \\chandra\\ for 44~ks during its first month of calibration-phase operations. The Galactic column density along this line of sight is $(3.9\\pm 0.4)\\times 10^{20}$~cm$^{-2}$ (Stark et~al. 1992), corresponding to an optical depth of $\\tau<0.02$ for the 2--8~keV band of primary interest here. In this paper we assume $H_0=70$~km~s$^{-1}$ Mpc$^{-1}$ and $q_0=\\frac{1}{2}$. ",
        "conclusions": "Our results, obtained from a small but well-defined sample of 2--8~keV sources, extend previous X-ray background studies in several ways. First, we have detected and securely identified hard X-ray sources about an order of magnitude fainter than has previously been possible. At our 2--8~keV flux limit of $\\approx$~(4--8)$\\times 10^{-15}$~erg~cm$^{-2}$~s$^{-1}$ (spatially dependent; see \\S2.2), we find ten sources (including AG1) in our 63.5~arcmin$^2$ search area, corresponding to a 2--8~keV source density of $570\\pm 180$~deg$^{-2}$. Even allowing for the possibility of one spurious source detection, this source density is still $\\approx 10$ times larger than previous number counts in this energy band (e.g., Ogasaka et~al. 1998; Giommi et~al. 1998), and down to $\\approx 2\\times 10^{-14}$~erg~cm$^{-2}$~s$^{-1}$ our number counts appear consistent with the \\asca\\ fluctuation analyses of Gendreau et~al. (1998). However, our source counts may be somewhat enhanced due to the presence of the cluster \\crsscluster\\ (see below). The fact that we detect sources down to $\\approx 4\\times 10^{-15}$~erg~cm$^{-2}$~s$^{-1}$ suggests that the number counts versus flux relation departs from the Euclidean form (the X-ray background would be resolved at $\\approx 10^{-14}$~erg~cm$^{-2}$~s$^{-1}$ without a break in the $\\log N$--$\\log S$ slope), although further data are clearly needed to quantify the break parameters.  Four of our five S3 sources with optical spectroscopy have $z<0.3$, which clearly exclude them from being associated with the cluster \\crsscluster\\ at $z=0.516$. The fifth, the QSO \\crssqso, differs in redshift from the cluster by $\\Delta z=0.023$ corresponding to a line-of-sight separation of $\\approx 100$~Mpc. While this QSO is certainly not a bound member of the cluster, it could be associated with the large-scale cosmic structure producing the cluster. As discussed in \\S3.3, if any of our blank-field sources (or AG1) lie in the cluster, they must be sub-$L^\\star$ galaxies producing large hard X-ray luminosities of $L_{2-8}\\simgt 4\\times 10^{42}$~erg~s$^{-1}$. Even in the most conservative (and unlikely) case, where we allow all objects with unknown $z$ to lie in the cluster, our cluster-corrected source density is still a factor of $\\approx 4.5$ times higher than previously attained by \\asca\\ and \\sax. The same statement obtains for possible gravitational lensing effects by the cluster.  We detect nine of our ten 2--8~keV sources in the 0.2--2~keV band. While our statistics are admittedly limited, this result is consistent with the finding by Giommi, Fiore \\& Perri (1998) that most hard X-ray sources have soft X-ray counterparts, and it extends this result downward in flux by about an order of magnitude (see \\S1 for discussion). Down to our flux limit, we can show with $>90$\\% confidence that hard-band only sources comprise $<40$\\% of the total hard-band source population. Deeper \\chandra\\ observations are needed to determine if a large population of hard-band only sources emerges at still fainter flux levels."
    },
    "0002/astro-ph0002317_arXiv.txt": {
        "abstract": "Analytic solutions for the evolution of cosmological linear density perturbations in the baryonic gas and collisionless dark matter are derived. The solutions are expressed in a closed form in terms of elementary functions, for arbitrary baryonic mass fraction.  They are obtained assuming $\\Omega=1$ and a time independent comoving Jeans wavenumber, $k_J$.  By working with a time variable $\\tau\\equiv \\ln(t^{2/3})$, the evolution of the perturbations is described by linear differential equations with constant coefficients. The new equations are then solved by means of Laplace transformation assuming that the gas and dark matter trace the same density field before a sudden heating epoch.  In a dark matter dominated Universe, the ratio of baryonic to dark matter density perturbation decays with time roughly like $\\exp(-5\\tau/4)\\propto t^{-5/6}$ to the limiting value $1/[1+(k/k_J)^2]$.  For wavenumbers $k>k_J/\\sqrt{24}$, the decay is accompanied with oscillations of a period $ 8\\pi/\\sqrt{24 (k/k_J)^2 -1}$ in $\\tau$.  In comparison, as $\\tau $ increases in a baryonic matter dominated Universe, the ratio approaches $1-(k/k_J)^2$ for $k\\le k_J$, and zero otherwise. ",
        "introduction": "Let $\\dx(t,k)$ and $\\db(t,k)$ be , respectively, the Fourier modes of   baryonic and dark matter density fluctuations. Let also $f_X$ and $f_b=1-f_x$ be the the mean mass fractions of  these two types of matter. We will restrict the analysis to perturbations in a flat universe without a cosmological constant. The linear equations governing the evolution of $\\db$ and $\\dx$ are (e.g., Bi et. al. 1992, Padmanabhan 1993, Gnedin \\& Hui 1998), \\begin{eqnarray} \\frac{\\dd^2 \\dx}{\\dd t^2} +2H\\frac{\\dd \\dx}{\\dd t}&=& \\frac{3}{2}H^2\\left(f_x \\dx + f_b \\db \\right) \\cr \\frac{\\dd^2 \\db}{\\dd t^2} +{2}H\\frac{\\dd \\db}{\\dd t}&=& \\frac{3}{2}H^2\\left(f_x \\dx + f_b \\db \\right) - \\frac{3}{2}H^2 \\left(\\frac{k}{k_J}\\right)^2  \\db \\; , \\label{lint} \\end{eqnarray} where $a(t)\\propto t^{2/3}$ is the scale factor, $H(t)=2/(3t)$ the Hubble function, and $k_J$ is the comoving Jeans wavenumber related to the speed of sound $c_s$ and the mean total (baryons plus dark) density, $\\bar \\rho=3H^2/(8\\pi G)$, by \\begin{equation} k_J=\\frac{a}{c_s}\\sqrt{4\\pi G \\bar \\rho}=\\sqrt{\\frac{3}{2}}\\frac{a H}{c_s} . \\end{equation}   For $f_x$ close to unity the gravity of the baryons is negligible and $\\dx \\propto a$. If also $k_J=const$ (i.e., $c_s^2\\propto 1/a$) and we impose, at some initial time $t_i$, the condition $\\db(t_i)=\\dx(t_i)/(1+(k/k_J)^2)$, then $\\db(t) =\\dx(t)/(1+(k/k_J)^2)$ at any later time, $t$.  In the next section we will show how to solve these equations with $0<le f_x\\le $ assuming that $\\db(t)=\\dx(t)$ for $t\\le t_i$, which is appropriate for a sudden reionization of the IGM at $t_i$.  In order to solve these equations we work with a new time variable $\\rp \\equiv \\ln (a)$ instead of $t$ (Nusser \\& Colberg 1998). In terms of $\\tau$ the equations  become, \\begin{eqnarray} \\frac{\\dd^2 \\dx}{\\dd \\rp^2} +\\frac{1}{2}\\frac{\\dd \\dx}{\\dd \\rp}&=& \\frac{3}{2}\\left(f_x \\dx + f_b \\db \\right) \\cr \\frac{\\dd^2 \\db}{\\dd \\rp^2} +\\frac{1}{2}\\frac{\\dd \\db}{\\dd \\rp}&=& \\frac{3}{2}\\left(f_x \\dx + f_b \\db \\right) -\\frac{3}{2} \\kappa^2 \\db \\; , \\label{lintau} \\end{eqnarray} where have defined  $\\kappa=k/k_J$.  When $\\kappa$ is constant, the differential equations (\\ref{lintau}) are linear  with constant coefficients and they can be solved by means of Laplace transformation. Since Laplace transforms are seldom used in cosmological studies, it seems prudent to briefly review their basic properties which are relevant to us. We refer the reader to Arfkin (1985) and references therein for mathematical details. The Laplace transform, $f(s)$, of a function $F(t)$, where $t\\ge 0$, is defined as \\begin{equation} f(s)\\equiv {\\cal L}\\{F(t)\\}=\\int_0^\\infty \\exp\\left(-st\\right) F(t) \\dd t . \\label{laplace:def} \\end{equation} We will need the Laplace transforms of first and second derivatives of a function.  Using (\\ref{laplace:def}) these transforms can be related to the $f(s)$ by \\begin{eqnarray} {\\cal L} \\{ F'(t)\\} &=& s f(s)- F(0) \\cr {\\cal L} \\{ F''(t)\\} &=& s^2 f(s)- s F(0) - F'(0) \\; , \\label{laplace:derv} \\end{eqnarray} where the prime and double prime denote first and second order derivatives, respectively. The Bromwich integral expresses  $F(t)$ in terms of $f(s)$ as \\begin{equation} F(t)=\\frac{1}{2\\pi i}\\int_{\\gamma -i\\infty}^{\\gamma+i\\infty} \\exp\\left(s t\\right) f(s)\\dd s \\label{bowich} \\end{equation} where $i=\\sqrt{-1}$ and $\\gamma$ is a real number chosen so that all  poles of $f(s)$ lie, in the complex plane, to the left of the vertical line defining the integration path. Therefore, by the residue theorem we have \\begin{equation} F(t)=\\sum \\left[ {\\rm residues} \\; {\\rm of} \\; \\exp\\left(s t\\right)f(s) \\right] \\label{resid} \\end{equation} As an example consider $f(s)=1/(s-s_1)(s-s_2)$ which has two simple poles at $s=s_1 $ and $s_2$. The residues of $\\exp({s t}) f(s)$ at these poles are $\\exp({s_1 t}) /(s_1-s_2)$ and $-\\exp({s_2 t}) /(s_1-s_2)$ so that, by (\\ref{resid}), $F(t)=[\\exp({s_1 t})- \\exp({s_2 t})]/(s_1-s_2)$. If $s_1=s_2$,  the function has a pole of order two at $s_1$. The residue in this case is $\\dd[(s-s_1)^2 \\exp({st})f(s))]/\\dd s$ evaluated at $s=s_1$. Therefore,  $F(t)=t \\exp({s_1 t})$. ",
        "conclusions": "We have found analytic solutions to the linear equations governing the evolution of baryonic and dark matter under four assumptions. First, the Universe is flat without cosmological constant.  Second, sudden reionization of the IGM. Third, the temperature of the low density IGM drops like $1/a$  so that the comoving Jeans length is time independent. Fourth, before reionization the IGM is cold and the baryonic and dark matter trace the same density and velocity fields.  Of these assumptions, only the fourth has a physical basis, at least before any heating has occurred and when the IGM temperature is low. This is also the only assumption which if changed, the equations can still be readily solved by Laplace transformation.  Unfortunately, relaxing any of the other assumptions complicates the analytic treatment of equations (\\ref{lintau}) by means of Laplace transformation.  For example, suppose that the Jeans wavenumber changes with time according to $a^\\beta$. Then the Laplace transformation of the term involving $k_J$ will yield $\\ldb$ at $(s+\\beta)$ while other terms involve $\\ldb(s)$.  Yet the solutions can be useful for semi-analytic modeling of the IGM. They offer a convenient improvement over the commonly used filter $1/[1+(k/k_J)^2]$ for generating gas fluctuations associated with a given mass density field.  The analytic solutions presented here were verified by a comparison with the solutions obtained by numerical integration of equations (\\ref{lint}).  All numerical and analytic solutions agreed up to the numerical accuracy."
    },
    "0002/astro-ph0002471_arXiv.txt": {
        "abstract": "The Cryogenic Dark Matter Search (CDMS) employs Ge and Si detectors to search for WIMPs via their elastic-scattering interactions with nuclei while discriminating against interactions of background particles.  CDMS data give limits on the spin-independent WIMP-nucleon elastic-scattering cross-section that exclude unexplored parameter space above 10~GeV~c$^{-2}$ WIMP mass and, at $> 84$\\% CL, the entire 3$\\sigma$ allowed region for the WIMP signal reported by the DAMA experiment. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002523_arXiv.txt": {
        "abstract": "We present the first results of a systematic analysis of radially truncated exponential discs for four galaxies of a complete sample of disc-dominated edge-on spiral galaxies.  \\\\ The discs of our sample galaxies are truncated at similar radii on either side of their centres.  With possible the exception of the disc of ESO 416-G25, it appears that the truncations in our sample galaxies are closely symmetric, in terms of both their sharpness and the truncation length.  However, the truncations occur over a larger region and not as abruptly as found in previous studies.  \\\\ We show that the truncated luminosity distributions of our sample galaxies, if also present in the mass distributions, comfortably meet the requirements for longevity.  The formation and maintenance of disc truncations are likely closely related to stability requirements for galactic discs. ",
        "introduction": "\\label{exptrunc.sec}  It is well-known from Freeman's (1970) work that the radial light distribution of the stellar component of high surface brightness galactic discs can be approximated by an exponential law of the form \\begin{equation} \\label{freeman.eq} L(R) = L_0 \\exp (- R / h_R) \\end{equation} where $L_0$ is the luminosity density in the galactic centre, {\\it R} the galactocentric distance and $h_R$ the disc scalelength.  However, for a few prominent high surface-brightness edge-on galaxies, it has initially been shown by van der Kruit \\& Searle (1981a,b, 1982a,b, hereinafter KS1--4), that at some radius $R_{\\rm max}$ (the truncation or cut-off radius of the galactic disc) the stellar luminosity distribution disappears asymptotically into the background noise (see also Jensen \\& Thuan 1982, Sasaki 1987, Morrison, Boroson \\& Harding 1994, Bottema 1995, Lequeux, Dantel-Fort \\& Fort 1995, N\\\"aslund \\& J\\\"ors\\\"ater 1997, Fry et al.  1999, Pohlen et al.  2000a,b).  In fact, the truncation of galactic discs does not occur instantly but over a small region, where the luminosity decrease becomes much steeper, having exponential scalelengths of order or less than a kiloparsec, opposed to several kpc in the exponential disc part (e.g., KS1--4, Jensen \\& Thuan 1982, Sasaki 1987, Abe et al.  1999, Fry et al.  1999).  An independent approach to obtain the statistics of truncated galactic discs, using a sample of galaxies selected in a uniform way, is needed in order to better understand the overall properties and physical implications of this feature.  In this paper we present the first results of a systematic analysis of disc truncations for a pilot sample of four ``normal'' spiral galaxies, drawn from the statistically complete sample of edge-on disc-dominated galaxies of de Grijs (1998). In Kregel, van der Kruit \\& de Grijs (2001) we will present a re-analysis of the global disc structures of the entire de Grijs (1998) sample, including a systematic analysis of the occurrence of truncated galactic discs.  Edge-on galaxies are particularly useful for the study of truncated galactic discs: since these disc cut-offs usually occur at very low surface brightness levels, they are more readily detected in highly inclined galaxies, where we can follow the light distributions out to larger radii.  In Sect.  \\ref{technical.sect} we outline the sample selection and our approach and methodology.  A detailed error discussion for our pilot sample is given in Sect. \\ref{results.sect}.  The most important science driver for the study of truncated galactic discs, discussed in Sect.  \\ref{dynamics.sect}, is that if the truncations seen in the stellar light are also present in the mass distribution, they would have important dynamical consequences at the disc's outer edges. ",
        "conclusions": "In this paper we have presented the first results of a systematic analysis of galactic disc structure in general and of radially truncated exponential discs in particular for a pilot sample of four ``normal'' disc-dominated edge-on spiral galaxies.  We have carefully considered the importance of (residual) dust, deviations from 90$^\\circ$ inclinations, and spiral arms, and concluded that these effects do not affect our results significantly.  We have also shown that the truncated discs in our sample galaxies are not caused artificially by inaccurate sky subtraction, but are real deviations from the radial exponential light profiles.  An independent approach to obtain the statistics of truncated galactic discs, using a sample of galaxies selected in a uniform way, is needed in order to better understand their overall properties and physical implications.  If the truncations seen in the stellar light are also present in the mass distribution, they would have important dynamical consequences at the disc's outer edges.  We have shown that the truncated luminosity distributions of our pilot sample galaxies, if also present in the mass distributions, comfortably meet the requirements for longevity.  The truncation radii, expressed in units of $h_R$, for the discs of ESO 201-G22, ESO 416-G25, and ESO 446-G18 are comparable to those found by KS1--4 and Bottema (1995), while ESO 446-G44 is truncated at much smaller radii.  In fact, the truncations of the discs of ESO 416-G25 and ESO 446-G44 are within the range found by Pohlen et al.  (2000a) for their sample of 31 nearby edge-on spiral galaxies.  In general, the discs of our sample galaxies are truncated at similar radii on either side of their centres, within the observational uncertainties, with the exception of ESO 201-G22.  With the possible exception of the disc of ESO 416-G25, it appears that our sample galaxies are fairly symmetric, in terms of both the sharpness of their disc truncations and the truncation length, although the truncations occur over a larger region and not as abruptly as found in previous studies.  The northern edge of the disc of ESO 416-G25 is very sharply truncated compared to its southern edge.  We believe that this may be explained by the fact that we likely observe the outer stellar envelope of a spiral arm, whereas on the southern side we are looking into the inside of a spiral arm."
    },
    "0002/astro-ph0002009_arXiv.txt": {
        "abstract": "We present  new VLA  and Effelsberg observations  of the  radio galaxy 8C\\,0821+695. We have obtained  detailed images in total intensity and polarization  of this 2  Mpc sized  giant.  The  magnetic field  has a configuration  predominantly parallel  to  the source  main axis.   We observe Faraday rotation at low frequencies, most probably produced by an ionized  medium external  to the radio  source. The  spectral index distribution is  that typical of  FR II radio galaxies,  with spectral indices  gradually steepening  from  the source  extremes towards  the core. Modeling  the spectrum in  the lobes using  standard synchrotron loss  models  yields the  spectral  age of  the  source  and the  mean velocity  of the  jet-head with  respect  to the  lobe material.   The existence of a  possible backflow in the lobe  is considered to relate spectral  with dynamical determinations  of the  age and  the velocity with respect to  the external medium. Through a  very simple model, we obtain a physical characterization of the jets and the external medium in which  the radio galaxy  expands.  The results in  8C\\,0821+695 are consistent with  a relativistic jet nourishing the  lobes which expand in a  hot, low  density halo.  We infer a  deceleration of  the source expansion velocity which we  explain through a progressive increase in the hot-spot size. ",
        "introduction": "Giant  Radio Galaxies  (GRGs)  constitute an  unusual  class of  radio sources  with projected  linear  sizes larger  than 1  Mpc\\footnote{We assume  H$_{0}$=75  $  \\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$ and  q$_{0}=0.5$ throughout this paper.}. These objects, although giant in size, do not stand out in luminosity, most  of them being of low surface-brightness (Subrahmanyan et al.  \\cite{subra2}).   Moreover, the enormous size of GRGs sometimes hampers the identification of emission from two distant lobes as  part of  an individual object  (e.g.  B2~1358+305;  Parma et al. \\cite{parma}).  Probably, these are  the reasons why to date there are only about 50 known GRGs (Ishwara-Chandra \\& Saikia \\cite{ishwara}; Schoenmakers et al. \\cite{arno1}).  Only recently, radio surveys like the Northern VLA Sky Survey (NVSS; Condon et al. \\cite{nvss}) or the Westerbork Northern Sky Survey (WENSS; Rengelink et al.   \\cite{wenss}) allow sensitive searches of  GRGs with adequate angular resolution.  Two  basic scenarios have  been envisaged  to explain  the outstanding sizes of GRGs.  First, their  lobes could be fed by extremely powerful central engines which  would endow the jets with  the necessary thrust to bore their long way  through the ambient medium. Second, GRGs could be  normal radio  sources  evolving in  very low-density  environments offering little  resistance to the  expansion of the jets.   While the first  possibility  requires  the  existence of  prominent  cores  and hot-spots, which are not always observed (Ishwara-Chandra \\& Saikia \\cite{ishwara}), the second possibility is supported by the high degree of polarization found in GRGs also at low frequencies (Willis \\& O'Dea \\cite{willis}) and  seems to  be the most  plausible scenario  in most cases (Mack et al.  \\cite{mack97b}).  It seems that the expansion of a radio galaxy in a low density environment during a time long enough to allow reaching Mpc  sizes, rather than higher than  usual radio powers or expansion velocities, are the two basic ingredients to build up the GRG population (Schoenmakers et al. \\cite{arno2}).  GRGs are located in regions hardly accessible via direct observations: they  do   not  reside  in  rich  galaxy   clusters  (Subrahmanyan  et al. \\cite{subra2})  and the X-ray emission around  their host galaxies is usually  weak (Mack et al.   \\cite{mack97a}).  However, information about  the  ambient medium  at  very  large  distances from  the  host galaxies  can  still  be  gained  through the  study  of  their  radio properties.  Most GRGs probe  the intergalactic medium (IGM) providing information about the density of matter outside galactic halos, adding important  observational constraints  to  current cosmological  models (Begelman \\& Cioffi \\cite{begelman}; Nath \\cite{nath}).  We discuss in  this paper new observations of  the GRG 8C\\,0821+695, a Fanaroff-Riley  type II  (Fanaroff \\&  Riley \\cite{fanaroff}).   It is optically  identified  with  a  faint  $M_{R}\\sim 22.2$  galaxy  at  a redshift  of 0.538  (Lacy et  al.  \\cite{lacy}).   Its  high redshift, compared to other GRGs, renders 8C\\,0821+695 a very interesting object since  it provides information  on the  external environment  at large cosmological  distances.  No  X-ray source  coincident with  the radio source is  found in the Bright  Point Source catalogue  from the ROSAT All  Sky  Survey.  At  the  distance  of  8C\\,0821+695, one  arcsecond corresponds to 4.9 kpc. ",
        "conclusions": "We  present new radio  observations made  with the  VLA and  the 100-m Effelsberg  radio  telescope, of  the  GRG  8C\\,0821+695 at  different frequencies  and  angular resolutions.  Our  data  have been  analyzed together with survey and literature data in order to study the details of  a high-redshift  GRG, and  obtain information  about  the external medium surrounding the radio source.  8C\\,0821+695  is a  straight 2  Mpc long  FR II  radio source,  with a north-to-south  arm-ratio of  0.89.  The  N-lobe contains  a hot-spot, responsible  for most  of the  source polarized  emission.  The S-lobe presents a  more relaxed structure,  without a well  defined hot-spot. At  high   frequencies  ($\\nu\\geq  1400$~MHz)   8C\\,0821+695  shows  a prominent  compact core.  We  have not  found  any trace  of the  jets nourishing the lobes.  The  spectral index  distribution over  the lobes  of  8C\\,0821+685 is typical  of FR II-type  radio galaxies,  showing a  gradual steepening from the outer ends towards the core.  The mean lobe spectral index is $\\alpha=1.1$.  The core has a flat spectrum with $\\alpha= 0.3$.  Using the  available data,  we have  made a  spectral-aging analysis  of the source lobes, providing the dependence of the spectral break frequency and the synchrotron  age with the distance from the  core. We obtain a mean  expansion velocity  of the  jet head  with respect  to  the lobe material of $0.08$~c, and a spectral  age of 42 Myr. This age determination might be affected by the possible existence of backflow of material in the lobe, being a lower limit to the true age. The age we derive for 8C\\,0821+695 is  of the order  of ages  estimated for  other GRGs (Schoenmakers et al. \\cite{arno2}).  We have studied  the $RM$ over the lobes  of 8C\\,0821+695, obtaining a smooth and uniform distribution,  which we ascribe to Faraday rotation mostly produced by the Galactic medium.  Equipartition conditions have been assumed  in order to derive  physical parameters of  the lobes at different  positions, yielding magnetic  fields, pressures  and energy densities that are consistent with estimated values in other GRGs.  Under very simple assumptions we have estimated physical parameters of the jet  and the  external medium of  8C\\,0821+695.  We find  that the present  expansion  velocity  is  significantly lower  than  the  mean expansion velocity even if backflow is allowed, implying the existence of deceleration.  We  explain this deceleration by an  increase of the cross-sectional  area of  the bow  shock at  the end  of the  jet with time. We  note that external  density estimates in the  literature for other GRGs usually  consider the contact surface measured  at the time of the observations together  with mean quantities (like the expansion velocity  derived   from  aging  arguments),   resulting  in  external densities lower than the density  we obtain.  However, our results are still consistent with GRGs evolving in poor density regions."
    },
    "0002/astro-ph0002379_arXiv.txt": {
        "abstract": "We present the first $UBV$ and $uvby$ photometric observations for the short period variable star HV~Ursae~Majoris classified as a field RRc variable. The observed differences between the consecutive minima and the lack of colour variations disagree with the RRc-classification and suggest the possible binary nature of HV~UMa. In order to reveal the real physical status of this star, we took medium resolution ($\\lambda/\\Delta \\lambda \\approx 11000$) spectra in the red spectral region centered at 6600 \\AA. Spectra obtained around the assumed quadratures clearly showed the presence of the secondary component.  An improved ephemeris calculated using our and Hipparcos epoch photometry is Hel. JD$_{\\rm min}=2451346.743\\pm0.001$, P$=0\\fd7107523(3)$. A radial velocity curve was determined by modelling the cores of H$\\alpha$ profiles with two Gaussian components. This approximative approach gave a spectroscopic mass ratio of q$_{\\rm sp}$=0.19$\\pm$0.03. A modified Lucy model containing a temperature excess of the secondary was fitted to the V light curve. The obtained set of physical parameters together with the parallax measurement indicate that this binary lies far from the galactic plane, and the primary component is an evolved object, probably a subgiant or giant star. The large temperature excess of the secondary may suggest a poor thermal contact between the components due to a relatively recent formation of this contact system. ",
        "introduction": "The first note on the possible light variability of HV~Ursae~Majoris (= HD~103576 = HIP~58157, $\\langle V \\rangle=8.69$, $\\Delta V=0.28$, $P=0\\fd355385$, $\\pi_{\\rm Hipp}=3.12\\pm1.23$ mas, $d_{\\rm Hipp}=320^{+210}_{-90}$ pc) was published by Penston (1973) who gave an 'uncertain' mark to the range of V-magnitude (``var? V=8.60--8.83''). The periodic nature of the light variation was discovered by the Hipparcos satellite (ESA 1997) and the star was classified as an RRc variable. There is a note in ESA (1997) about the possibility of a double period but no firm conclusion was drawn. ESA (1997) gives a spectral type A3, while Slettebak \\& Stock (1959) published A7.  We started a long-term observational project of Str\\\"omgren photometry and spectroscopy of the newly discovered bright Hipparcos variables. The first results have already appeared in Kiss et al. (1999a, b). Since the period, spectral type and light curve do not exclude the possibility of wrong classification, accurate determination of the fundamental physical parameters is highly desirable.  The main aim of this paper is to present the first $UBV$ and $uvby$ photometry for HV~UMa. Also, our radial velocity measurements are the first time-resolved spectroscopic observations of this star to date. The paper is organised as follows: the observations are described in Sect.\\ 2, Sect.\\ 3 deals with data analysis and the obtained physical parameters, while a discussion of the results is given in Sect.\\ 4. A final list of conclusions is presented in Sect.\\ 5. ",
        "conclusions": "The new results presented in this paper can be summarized as follows.  \\noindent 1. We reported the first $uvby$ photometric and medium resolution spectroscopic observations of the short-period variable star HV~Ursae~Majoris. Contrary to the RRc classification given by ESA (1997), the star turned out to be a new contact binary with early-type components.  \\noindent 2. An improved ephemeris was determined using our and Hipparcos epoch photometric data: Hel.JD$_{min}$=2451346.743+0.7107523(3)$\\cdot$E. There is no indication of changing period over almost 10 years.  \\noindent 3. A radial velocity curve was measured directly from the H$\\alpha$ profiles by two-component Gaussian fitting of the line core regions in spectra recorded around the quadratures. We calculated a spectroscopic mass ratio of 0.19$\\pm$0.03.  \\noindent 4. The effective temperature and the surface gravity were determined using the mean Str\\\"omgren indices, synthetic colours of Kurucz (1993) and theoretical (ATLAS9) line profiles giving $T_{eff}=7000\\pm200$ K and log~$g=4.0\\pm0.3$. The light curve modelling resulted in a complete set of physical parameters.  \\noindent 5. The physical parameters of HV UMa together with the parallax measurement indicate that this binary is situated far from the galactic plane, and the primary component is an evolved object, probably a subgiant or giant star. The large temperature excess of the secondary may be indicative of a poor thermal contact between the components due to a relatively recent formation of this contact system via case B mass transfer."
    },
    "0002/astro-ph0002186_arXiv.txt": {
        "abstract": "Images of neutral Hydrogen 21cm absorption and radio continuum emission at 1.4 GHz from Mrk 273 were made using the Very Long Baseline Array and Very Large Array. These images reveal a gas disk associated with the northern nuclear region with a diameter of 0.5$''$ (370 pc), at an inclination angle of 53$^o$.  The radio continuum emission is composed of a diffuse component plus a number of compact sources.  This morphology resembles those of nearby, lower luminosity starburst galaxies.  These images provide strong support for the hypothesis that the luminosity of the northern source is dominated by an extreme compact starburst.  The HI 21cm absorption shows an east-west gradient in velocity of 450 km s$^{-1}$ across 0.3$''$ (220 pc), implying an enclosed mass of 2$\\times$10$^{9}$ M$_\\odot$, comparable to the molecular gas mass.  The brightest of the compact sources may indicate radio emission from an active nucleus (AGN), but this source contributes only 3.8$\\%$ to the total flux density of the northern nuclear region.  The HI 21cm absorption toward the southeast radio nucleus suggests infall  at 200 km s$^{-1}$ on scales $\\le$ 40 pc, and the southwest near IR  nucleus is not detected in high resolution radio continuum images. ",
        "introduction": "Observations of Mrk 273 were made on May 31 and June 6, 1999 with the Very Long Baseline Array (VLBA), including the phased Very Large Array (VLA) as an element in the very long baseline array. The pass band was centered at the frequency of the neutral hydrogen 21cm line at a heliocentric redshift of: z =  0.0377, or cz = 11300 km s$^{-1}$. The total bandwidth was 16 MHz, using two orthogonal polarizations, 256 spectral channels, and 2 bit correlation. The total on-source observing  time was 13.4 hrs.  Data reduction was performed using the Astronomical Image Processing System (AIPS) and AIPS++.  Standard {\\sl a priori} gain calibration was performed using the measured gains and system temperatures of each antenna.  The compact radio source J1337+550 was observed every 5 minutes, and this source was used to determine the initial fringe rates and delays.  The source 3C 345 was used to calibrate the frequency dependent gains (band pass calibration). The source J1400+621 was used to check the absolute gain calibration. The results showed agreement of observed and expected flux densities to within 3$\\%$.  After application of the delay and rate solutions, and band pass calibration, a continuum data set for Mrk 273 was generated by averaging off-line channels. This continuum data set was then used for the hybrid imaging process, which involves iterative imaging and self-calibration of the antenna-based complex gains (Walker 1985). The final iteration involved both phase and amplitude calibration with a 3 minute averaging time for phases and 15 minutes for amplitudes.  The self-calibration solutions were applied to the spectral line data set. The spectral line data were then analyzed at various spatial and spectral resolutions by tapering the visibility data, and by smoothing in frequency.  The continuum emission was subtracted from the spectral line visibility data using UVLIN.  Images of the line and continuum data were deconvolved using the Clark `CLEAN' algorithm as implemented in IMAGR.  For the radio continuum images we also employed the multi-resolution CLEAN algorithm as implemented in AIPS++ (Holdaway and Cornwell 1999).  Results were consistent for all image reconstruction algorithms, and we present the naturally weighted Clark CLEAN continuum images in the analysis below. The full resolution of the naturally weighted images is 10 mas. We also present images at 50 mas resolution made using a Gaussian taper of the visibilities. ",
        "conclusions": ""
    },
    "0002/astro-ph0002465_arXiv.txt": {
        "abstract": "HD~220392 (HR~8895), the brightest member of the visual double star CCDM~23239-5349, is a new short-period variable bright star, probably of the $\\delta$ Scuti type. The period analysis performed on the complete set of definitive Geneva photometry as well as on the data obtained at the ESO 0.5m telescope shows two periodicities of about 4.7 and 5.5 cycles per day (cpd) with amplitudes of 0.014 and 0.011 mag respectively. A similar period search on the (smaller) dataset obtained for the 1 mag fainter B-component, HD~220391, however shows no periodicity with an amplitude significantly above the noise level of the data (about 0.006~mag). This difference in variability behaviour is discussed from the consideration that {\\it both} stars form a common origin pair and are located in the $\\delta$ Scuti instability strip. ",
        "introduction": "The very wide double star CCDM~23239-5349 is an interesting study case of a pulsating star within a common origin pair or wide binary. The detailed investigation of the difference in variability and physical parameters between two components of a physical couple is particularly worthwhile when both stars are located in the same area of the colour-magnitude diagram, in this case {\\it both} components are in the $\\delta$ Scuti instability strip. The aim of such a study is to search for clues to understand what factors determine the pulsation characteristics such as modes and amplitudes among $\\delta$ Scuti stars in general.\\\\ The Hipparcos satellite measurements confirm what was already  hinted by the ground-based astrometric data in the Washington Double Star Catalogue (WDS 1996.0, Worley \\& Douglass \\cite{wor97}), namely that the wide angular separation of 26.5 \\arcsec of the system is accompanied by a very small relative proper motion ($\\Delta\\mu_{\\alpha^{\\ast }_{B-A}}\\simeq -2.44$~milli-arcsec/yr (mas/yr), $\\Delta \\mu_{\\delta_{B-A}}\\simeq +1.57$~mas/yr with errors of the same order). The new parallaxes are furthermore compatible to better than $1.5\\,\\sigma$. This may indicate a common origin if not a true physical association (Sect.~\\ref{sec:ass}).  Regular short-period light variations on a time scale of $\\simeq$ 5 hr have been detected for the brightest component of this visual double star (Lampens \\cite{lam92}). We describe the available observations and the reduction methods in Sect.~\\ref{sec:obs}. The results of the period analyses are presented in Sect.~\\ref{sec:res}. Also included is the analysis of a selection of the Hipparcos Epoch Photometry data. We discuss the nature of the association and of the variability in Sect.~\\ref{sec:nat}. Finally we draw our conclusions in Sect.~\\ref{sec:con} and we explain why additional observations for both stars of the system would be highly desirable. ",
        "conclusions": "\\label{sec:con}  Binary and multiple systems with pulsating variable components offer a unique opportunity of coupling the information obtained by astrometric means (association type - parallax - total mass) to the astrophysical quantities gained from the photometry /spectroscopy (luminosity ratio - colours - pulsation characteristics)(see Lampens \\& Boffin \\cite{lam00} for a review of $\\delta$ Scuti stars in stellar systems). The detailed investigation of the differences in variability and simul\\-ta\\-neous\\-ly in physical properties between two components of a binary system may provide clues with respect to the pulsation: differences in origin and age can be ruled out as well as differences in overall chemical composition. Stronger constraints exist for the determination of the position of the components in the H-R diagram, there is therefore less ambiguity in determining the evolutionary status and the mass than in the case of single variable stars. This is important when one of the components is located in the zone where evolutionary tracks are bent (e.g. near the end of the core hydrogen burning phase).  A relevant question is what factors determine the pulsation characteristics (the amplitudes and the modes) in the $\\delta$ Scuti instability strip? We addressed this from the point-of-view of two bright A/F-type stars that are both located in the $\\delta$ Scuti instability strip and that are shown to be physically associated, i.e. they either form a common origin pair or they are the components of a true wide binary system. In this case, evolution (and mass) is the most pronounced physical difference between both stars and it is very probable that this is the cause for the observed difference in variability behaviour. Further observations are needed, the more that, since there is no evidence for any metal lines in the spectra, a comprehensive variability analysis of this system might also help explaining the presence of non-variable, non-metallic stars in the instability strip.  In the light of the discussion by Solano \\& Fernley (\\cite{sol97}) on the relation between rotational velocity and amplitude, we noted the remarkable similarity of the projected rotational velocities: both stars are rather fast rotators. If fast rotation favours pulsation of the $\\delta$ Scuti type, we expect to find short-period variability for the B-component as well! Since it is less evolved than its brighter companion, smaller amplitudes are expected. This is another reason why intensive monitoring of this southern system is certainly worthwhile. In our example it was very easy to identify the short-period pulsating component and the information obtained from the astrometry could be coupled to the astrophysical parameters of each component individually. Even better would be to investigate these characteristics in a close visual binary for which information on the orbital motion can also be derived. This will allow to obtain a direct estimation of the stellar mass, independent from the choice of modelisation. The derivation of the pulsation constant will be more straightforward (the error on the mass defines the accuracy of Q). More cases like this one should be looked into (see Frandsen et al. \\cite{fra95}).  With this application in mind, we made a crossidentification between the Annex of Variable Stars and the Annex of Double and Multiple Stars from the Hipparcos Catalogue (ESA \\cite{esa97}). Some 2500 systems with at least one variable component have been identified. But the description of the variability or the light curve in the Annex always refer to the combined magnitudes. Additional observations should help identify which component is variable and which are the binaries that offer the opportunity of coupling the information obtained by astrometric means to the physical properties in order to obtain a consistent picture of the system and its components.\\newline  {\\bf Acknowledgements} We thank the Geneva team (especially Dr. G. Burki) for the telescope time put at our disposal in June~1990 and September~1991. Dr. Sperl is kindly acknowledged for making the programme Period98 available for this application. We appreciate the help of Dr. L. Eyer (Geneva Observatory) in the selection of the Hipparcos Epoch Photometry data. We thank our colleague, Dr. J. Cuypers, for a critical reading and the referee, Dr. P. North, for constructive comments on this manuscript. The Geneva data can be requested from the Geneva photometry team. The ESO data are available on request from the authors. This work has made use of the Simbad database, operated by the {\\it Centre de Donn\\'ees astronomiques de Strasbourg} (France). \\newline"
    },
    "0002/astro-ph0002303_arXiv.txt": {
        "abstract": "MACHe3 (MAtrix of Cells of superfluid \\hetro) is a project of a new detector for direct Dark Matter (DM) search. A cell of superfluid \\hetrois has been developed and the idea of using a large number of such cells in a high granularity detector is proposed. \\noindent This paper presents, after a brief description of the superfluid \\hetrois cell, the simulation of the response of different matrix configurations allowing to define an optimum design as a function of the number of cells and the volume of each cell. The background rejection, for several configurations, is presented both for neutrons and \\gams of various kinetic energies. ",
        "introduction": "As previously suggested \\cite{vieux1,vieux2}, superfluid \\hetrois provides a suitable working medium for the detection of low energy recoil interactions. Recent studies \\cite{prl95} have shown the possibility to use a superfluid \\hetrois cell at ultra low temperatures (T$\\simeq$100 $\\mu$K). The primary device consisted of a small copper cubic box (V$\\simeq$ 0.125 cm$^{3}$) filled with \\hetro. It is immersed in a larger volume containing liquid \\hetrois and thin plates of copper nuclear-cooling refrigerant, see fig. \\ref{fig:design}. Two vibrating wires are placed inside the cell, forming a Lancaster type bolometer\\cite{prl95}. A small hole on one of the box walls connects the box to the main \\hetrois volume, thus allowing the diffusion of the thermal excitations of the \\hetrois generated by the energy deposited in the bolometer by the interacting particle.\\\\ This high sensitivity device is used as follows : the incoming particle deposits an amount of energy in the cell, which is converted into \\hetrois quasiparticles. These are detected by their damping effect on the vibrating wire. It must be pointed out that the size of the hole governs the relaxing time (quasiparticles escape time) and the Q factor of the resonator governs the rising time, see figure \\ref{fig:signal}. The present device has a rather high Q factor (Q $\\simeq 10^{4}$), giving a rising time of the order of one second. Although the primary experiment was still rudimentary, it has allowed to detect signals down to a threshold of 1 keV \\cite{prl95}. Many ideas are under study to improve the sensitivity of such a cell. Recently, the fabrication of micromechanical silicon resonators has been reported \\cite{trique} and the possibility to use such wires at ultra-low temperatures is under study.\\\\ The aim of the present article is to show that, by using a large number of these cells,  a high granularity superfluid \\hetrois detector could be used for direct Dark Matter (DM) search\\footnote{It should be noticed that such a device need to be placed in an underground site to reduce cosmic rays (mainly muons) background. It should also be surrounded by neutron and \\gam shieldings.}. For this purpose we have evaluated, by simulation of different kind of background events, the rejection coefficients that may be achieved with such a device. ",
        "conclusions": "In this prospective paper, we have demonstrate that a large matrix ($\\sim$ 1000 cells) of large cells (125 $cm^{3}$) is the preferred design for a superfluid \\hetrois detector searching for DM, as far as background rejection is concerned. An experimental work needs to be done to demonstrate the possibility to use such a large volume of superfluid $ ^{3}{\\mathrm{He}}$-B at ultra-low temperatures. This work has evaluated the background rejection of a high granularity superfluid \\hetrois detector for a large range of kinetic energies, both for neutrons and $\\gamma$-rays. It has been shown that, by means of correlation among the cells and energy loss measurement, a high rejection may be obtained for $\\gamma$-ray, neutron and muon background. Using the measured muon and neutron flux in an underground laboratory, we have evaluated the contamination to be one order of magnitude (two orders for muons) less than the expected \\neut rate. For background rejection purpose, the main advantage of a superfluid \\hetrois detector is to present a high rejection against neutron background, mainly because of the high capture cross-section at low energy. As neutrons interact {\\it a priori} like \\neutt, they are the ultimate background noise for DM detectors.\\\\  \\textbf{Acknowledgements}\\\\ The authors are grateful to D. Kerdraon and L. Perrot for the help concerning the use of MCNP calculation code, and also F. Ohlsson-Malek for the fruitful discussions on the GEANT code."
    },
    "0002/astro-ph0002245_arXiv.txt": {
        "abstract": "A model of supernovae feedback during disc galaxy formation is developed. The model incorporates infall of cooling gas from a halo and outflow of hot gas from a multiphase interstellar medium (ISM). The star formation rate is determined by balancing the energy dissipated in collisions between cold gas clouds with that supplied by supernovae in a disc marginally unstable to axisymmetric instabilities.  Hot gas is created by thermal evaporation of cold gas clouds in supernovae remnants, and criteria are derived to estimate the characteristic temperature and density of the hot component and hence the net mass outflow rate. A number of refinements of the model are investigated, including a simple model of a galactic fountain, the response of the cold component to the pressure of the hot gas, pressure induced star formation and chemical evolution. The main conclusion of this paper is that low rates of star formation can expel a large fraction of the gas from a dwarf galaxy.  For example, a galaxy with circular speed $\\sim 50\\;\\kms$ can expel $\\sim 60$--$80\\%$ of its gas over a time-scale of $\\sim 1$ Gyr, with a star formation rate that never exceeds $\\sim 0.1 M_\\odot$/year. Effective feedback can therefore take place in a quiescent mode and does not require strong bursts of star formation. Even a large galaxy, such as the Milky Way, might have lost as much as $20\\%$ of its mass in a supernovae driven wind.  The models developed here suggest that dwarf galaxies at high redshifts will have low average star formation rates and may contain extended gaseous discs of largely unprocessed gas. Such extended gaseous discs might explain the numbers, metallicities and metallicity dispersions of damped Lyman alpha systems. \\vskip 0.2 truein ",
        "introduction": "Since the pioneering paper of White and Rees (1978), it has been clear that some type of feedback mechanism is required to explain the shape of the galaxy luminosity function in hierarchical clustering theories. The reason for this is easy to understand; if the power spectrum of mass fluctuations is approximated as a power law $P(k) \\propto k^n$, the Press-Schechter (1974) theory for the distribution of virialized haloes predicts a power law dependence at low masses \\begin{equation} {d N(m) \\over dm} \\propto m^{-(9-n)/6}. \\label{Int1} \\end{equation} For any reasonable value of the index $n$ ($n \\approx -2$ on the scales relevant to galaxy formation in cold dark matter (CDM) models), equation~(\\ref{Int1}) predicts a much steeper mass spectrum than the observed faint end slope of the galaxy luminosity function, $dN(L)/dL \\propto L^\\alpha$, with $\\alpha \\approx -1$ (Efstathiou, Ellis and Peterson 1988, Loveday \\etal$\\;$ 1992, Zucca \\etal$\\;$ 1997). Furthermore, the cooling times of collisionally ionised gas clouds forming at high redshift are short compared to the Hubble time (Rees and Ostriker 1977, White and Rees 1978). Thus, in the absence of feedback, one would expect that a large fraction of the baryons would have collapsed at high redshift into low mass dark matter haloes,  in contradiction with observations.  In reality, there are a number of complex physical mechanisms that can influence galaxy formation and these need to be understood if we are to construct a realistic model of galaxy formation.  In the `standard' cold dark matter model ({\\it i.e.} nearly scale invariant adiabatic perturbations), the first generation of collapsed objects will form in haloes with low virial temperatures ($T \\simlt 10^4$K, characteristic circular speeds $v_c \\simlt 20 \\kms$). Molecular hydrogen is the dominant coolant at such low temperatures and so an analysis of the formation of the first stellar objects requires an understanding of the molecular hydrogen abundance and how this is influenced by the ambient ultraviolet radiation field (Haiman, Rees and Loeb 1997, Haiman, Abel and Rees 1999).  As the background UV flux rises, the temperature of the intergalactic medium will rise to $\\sim 10^4$K ({\\it e.g.} Gnedin and Ostriker 1997) and the UV background will reduce the effectiveness of cooling in low density highly ionized gas (Efstathiou 1992). A UV background can therefore suppress the collapse of gas in regions of low overdensity. It is this low density photoionized gas that we believe accounts for the Ly$\\alpha$ absorption lines (Cen \\etal$\\;$ 1994, Hernquist \\etal$\\;$ 1996, Theuns \\etal$\\;$ 1998, Bryan \\etal$\\;$ 1999). Photoionization can also suppress the collapse of gas in haloes with circular speeds of up to $v_c \\sim 20$--$30\\kms$. However, numerical simulations have shown that a UV background cannot prevent the collapse of gas in haloes with higher circular speeds, though it can reduce significantly the efficiency with which low density gas is accreted onto massive galaxies (Quinn, Katz and Efstathiou 1996, Navarro and Steinmetz 1997).  To explain the galaxy luminosity function, feedback is required in galaxies with circular speeds $v_c \\simgt 50 \\;\\kms$ with characteristic virial temperatures of $\\simgt 10^5\\;$K. Energy injection from supernovae is probably the most plausible feedback mechanism for systems with such high virial temperatures. Winds from quasars might also disrupt galaxy formation (Silk and Rees 1998) or, more plausibly, limit the growth of the central black hole (Fabian 1999). Here we will be concerned exclusively with supernova driven feedback and will not consider feedback from an active nucleus. Simple parametric models of supernovae feedback were developed by White and Rees (1978) and White and Frenk (1991) and form a key ingredient of semi-analytic models of galaxy formation ({\\it e.g.} Kauffmann, White and Guiderdoni 1993, Lacey \\etal$\\;$ 1993, Cole \\etal$\\;$ 1994, Baugh \\etal$\\;$ 1996, 1998, Somerville and Primack 1999).  In this paper, we develop a more detailed model of the feedback process itself.  Previous papers on supernovae feedback include those of Larson (1974), Dekel and Silk (1986) and Babul and Rees (1992). These authors compute the energy injected by supernovae into a uniform interstellar medium (ISM) and apply a simple binding energy criterion to assess whether the ISM will be driven out of the galaxy. The feedback process in these models is explosive, operating on the characteristic timescale of $\\sim 10^6$ --$10^7$ yrs for supernova remnants to overlap.  This is much shorter than the typical infall timescale of hot gas in the halo, begging the question of how a resevoir of cold gas accumulated in the first place. The present paper differs in that we model the ISM as a two-phase medium consisting of cold clouds and a hot pressure confining medium, {\\it i.e.} as a simplified version of the three-phase model of the ISM developed by McKee and Ostriker (1977, hereafter MO77). The cold component contains most of the gaseous mass of the disc and is converted into a hot phase by thermal evaporation in expanding supernovae remnants. In this type of model, the cold phase can be lost gradually in a galactic wind as it is slowly converted into a hot phase.   The main result of this paper is that low rates of star formation can expel a large fraction of the baryonic mass in dwarf galaxies over a relatively long timescale of $\\sim 1$ Gyr. We therefore propose that effective feedback can operate in an steady, unspectacular mode; strong bursts of star formation and superwind-like phenomena ({\\it e.g.} Heckman, Armus and Miley 1990) are not required, although galaxies may experience additional feedback of this sort. In fact, hydrodynamic simulations suggest that nuclear starbursts are ineffective in removing the ISM from galaxies with gas masses $\\simgt 10^6 M_\\odot$ (Mac Low and Ferrara 1999, Strickland and Stevens 1999) because hot gas generated in the nuclear regions is a expelled in a bipolar outflow without coupling to the cool gas in the rest of the disc.  This result provides additional motivation for investigating a ``quiescent'' mode of feedback.  Silk (1997) describes a model which is similar, in some respects, to the model described here. However, the  model described here is  more detailed and allows a crude investigation of the radial properties of a disc galaxy during formation.   The layout of this paper is as follows. A simple model of star formation regulated by disc instabilities is described in Section 2. This is applied to `closed box' ({\\it i.e.}  no infall or outflow of gas) models of disc galaxies neglecting feedback. Section 3 describes a model of the interaction of expanding supernovae shells in a two-phase ISM. This section is based on the model of MO77, but instead of focussing on equilibrium solutions that might apply to our own Galaxy, we compute the net rate of conversion of cold gas to hot gas incorporating the model for self-regulating star formation. This yields the temperature and density of the hot phase as a function of time and radius within the disc.  Section 4 revisits the model of Section 2, but includes simultaneous infall and outflow of gas.  This model is extended in Section 5 to include a galactic fountain, the pressure response of the cold ISM to the hot phase, and a model of chemical evolution. Section 6 describes some results from this model and discusses the effects of varying some of the input parameters.  In addition, the efficiency of feedback is computed as a function of the circular speed of the surrounding dark matter halo. Our conclusions are summarized in Section 7. Although we focus on disc galaxies in this paper, a similar formalism could be applied to the formation of bulges if the assumption that gas conserves its angular momentum during collapse is relaxed. ",
        "conclusions": "The main aim of this paper has been to show that supernovae driven feedback can operate in a quiescent mode and that high rates of star formation are not required to drive efficient feedback. In dwarf galaxies feedback occurs on an infalling timescale and so can extend over a period of $\\sim 1$ Gyr. In the feedback model developed here, cold gas clouds are steadily evaporated in expanding supernovae remnants and converted into a hot component.  Critically, the rate at which cold gas is evaporated can exceed the rate at which mass is converted into stars. If the temperature of the hot component is high enough, a wind will form and the hot gas can escape from the halo (provided the interaction with infalling gas can be ignored). If the temperature of the hot component is not high enough for it to escape from the halo, it will cool and fall back down to the  disc in a galactic fountain. Some characteristic features of the models are as follows:  \\smallskip  \\noindent (i) In a Milky Way type system, feedback from supernovae may drive out some of the gas from the halo in the early phases of evolution ($t \\simlt 0.3$ Gyr) when the star formation rate is high and the temperature of the hot phase exceeds $\\sim 5 \\times 10^6$ K.  For plausible sets of parameters, perhaps $20$ -- $30 \\%$ of the final stellar mass might escape from the galaxy. At later times, the temperature of the hot phase drops to $T \\sim 10^6$ K and the evaporated gas cycles within the halo in a galactic fountain.  \\smallskip  \\noindent (ii) In a dwarf galaxy with a circular speed $\\sim 50 \\; \\kms$, expanding supernovae remnants can convert the cold interstellar medium efficiently into a hot component with a chacteristic temperature of a few times $10^5$ K. This evaporated gas can escape from the halo in a cool wind. Typically, only about $10\\%$ of the baryonic material forms stars. Gas  accreted from the halo at  $\\simgt 1$ Gyr forms an extended gaseous disc which, according to the self-regulating star formation model used here, can  survive for longer than a Hubble time without converting into stars.  \\smallskip  \\noindent (iii) The feedback model developed here is meant to provide a sketch of how feedback might operate in a multi-phase interstellar medium. The model contains a number of obvious over-simplifications. For example, we have neglected any interaction of the outflowing gas with the infalling medium, we have not addressed the origin of the cold cloud spectrum, ignored the dense molecular cloud component of the ISM and neglected any local dissipation of supernovae energy in star forming regions.  These effects, and other processes, are undoubtedly important in determining the efficiency of feedback.  Nevertheless, the simplified model presented here contains some interesting features. Firstly, the model shows how positive feedback (via pressure induced star formation) and negative feedback (via outflowing gas) can occur {\\it simultaneously}. Secondly, the models are remarkably insensitive to uncertain physical parameters, in particular, thermal conduction would need to be suppressed relative to its ideal value by factors of much more than $100$ to qualitatively change the model.  If thermal conduction is highly suppressed, it may be possible to construct a qualitatively similar model to the one presented here in which cold gas is converted into hot gas in shocks.  \\smallskip  \\noindent (iv) The self-regulated star formation and feedback models described here provide physically based models for the star formation timescale and feedback efficiency as a function of the parameters of the halo. The star formation timescale and feedback efficiency (or timescale) are taken as free parameters in semi-analytic models of galaxy formation ({\\it e.g.} Cole \\etal$\\;$ 1999, Kauffmann \\etal$\\;$ 1994) and are critically important in determining some of the key predicted properties of these models, for example, the faint end slope of the galaxy luminosity function and the star formation history at high redshifts (see {\\it e.g.}  Somerville and Primack 1999). It is therefore important that we develop a theoretical understanding of these parameters (as attempted crudely here) and also that ways are found to constrain these parameters observationally. The results presented here show that supernovae feedback is much less effective than assumed in some earlier semi-analytic models (Cole \\etal$\\;$ 1994, Baugh \\etal$\\;$ 1996) but is closer to the more gentler feedback prescriptions used in more recent models (Cole \\etal$\\;$ 1999, Somerville and Primack 1999).  The feedback model described in this paper has a number of consequences and raises some problems which are summarized below.  \\smallskip  \\noindent (i) {\\it Evidence for outflows:} According to the models described here, outflows with speeds of $\\sim 200\\; T_{h6}^{1/2}\\kms$ should be common in high redshift galaxies. There is evidence for an outflow of $\\sim 200 \\; \\kms$, a mass loss rate of $\\sim 60 M_\\odot/{\\rm yr}$ and a star formation rate of $\\sim 40 M_\\odot/{\\rm yr}$ in the gravitationally lensed Lyman break galaxy MS1512-cB58 (Pettini \\etal$\\;$ 2000). The outflow velocity in this galaxy is consistent with our models, but the star formation and mass loss rates (which are highly uncertain) are high. The most likely explanations are either that MS1512-cB58 is a massive galaxy driving an outflow that will remain bound to the system, or that it is a less massive system undergoing a burst of star formation. In addition to direct detection of outflowing gas, winds may have other observational consequences.  The winds from dwarf galaxies will cool rapidly (see Appendix B). Wang (1995b) has suggested that photoionized gas clouds formed in the cooling wind might contribute to the Ly$\\alpha$ forest. Nulsen, Barcon and Fabian (1998) suggest that outflows caused by bursts of star formation in dwarf galaxies might even produce damped Ly$\\alpha$ systems.   \\smallskip  \\noindent (ii) {\\it Damped Ly$\\alpha$ systems:} The extended gaseous discs that form around dwarf galaxies in our models have low metallicities because they have low rates of star formation. If this is correct, then this largely unprocessed gas would dominate the cross section for the formation of damped Ly$\\alpha$ absorbers. The metallicities of most of these systems would be low, but would show a large scatter because some lines of sight will pass close to the central regions of galaxies containing  gas of high metallicity. This is broadly in agreement with what is observed. Extended gaseous discs would be vulnerable in tidal interactions. Some of the gas might be stripped and some might be transported into the central regions to be  converted into stars and hot gas. The evolution of $\\Omega_{\\rm HI}$ determined from damped Ly$\\alpha$ systems ({\\it e.g.}  Storrie-Lombardi, McMahon and Irwin 1996) might have more to do with infall, feedback and tidal disruption than with the cosmic star formation history.  \\smallskip  \\noindent (iii) {\\it Angular momentum conservation:} In hydrodynamic simulations, gas is found to cool effectively in sub-units during the formation of a protogalaxy. These sub-units lose their orbital angular momentum to the halo as they spiral towards the centre and merge. Hence the gas does not conserve angular momentum during the formation of a massive galaxy.  (Navarro and Benz 1991, Navarro and Steinmetz 1997, Weil Eke and Efstathiou 1998, Navarro and Steinmetz 2000). In fact, in the absence of feedback it has proved impossible to form discs with angular momenta similar to those of real disc galaxies starting from CDM initial conditions.  In the models described here, it has been assumed for simplicity that the specific angular momentum of the gas is conserved during collapse. This assumption could easily be relaxed. However, the feedback model decribed here suggests that the numerical simulations miss some important physics. Firstly, it is preferentially the low angular momentum gas, infalling in the early stages of evolution, that is most likely to be ejected from a developing protogalaxy. Secondly, supernovae driven feedback may help to solve the angular momentum momentum problem by ejecting gas efficiently from sub-units. The ejected gas may then infall at later times when the halo is less sub-structured, approximately conserving its angular momentum (Weil \\etal 1998, Eke, Efstathiou and Wright 2000).    \\smallskip  \\noindent (iv) {\\it Implementing feedback in numerical simulations:} There have been a number of attempts to implement supernovae feedback in gas dynamical numerical simulations ({\\it e.g.} Katz 1992, Navarro and White 1993, Navarro and Steinmetz 2000). These involve either heating the gas around star forming regions (which is ineffective because the energy is quickly radiated away) or reversing the flow of infalling gas. An implementation of the feedback model described here is well beyond the capabilities of present numerical codes. It would require modelling several gas phases, a cold interstellar medium, a hot outflowing medium and an infalling component, including mass transfer between each phase. It might be worth attempting simpler simulations in which cold high density gas is added to the halo beyond the virial radius at a rate that is determined by the local star formation rate.   \\smallskip  \\noindent (vi) {\\it Starbursts vs quiescent feedback:} It is likely that starburts are more common at high redshift because of the increased frequency of galaxy interactions. Starbursts could contribute to supernovae driven feedback in addition to the quiescent mode described here. However, at any one time, our models suggest that the cold gas component will have a mass of only $20$ -- $50\\%$ of the the mass of the stellar disk. Even if a substantial fraction of this gas is transported towards the centre of a galaxy in a tidal encounter (see {\\it e.g.}  Barnes and Hernquist 1996) and is subsequently ejected in a superwind, this mode of feedback will be inefficient because the mass of gas involved is a small fraction of the total gas mass ejected in the quiescent feedback mode over the lifetime of the galaxy.   \\smallskip  \\noindent (vii) {\\it Metallicity ejection:} The mean metallicity of the gas ejected from a dwarf galaxy is typically about $Z_\\odot/10$ in our models, and comparable to the mean metallicity of the stars in the final galaxy. Yet typically a dwarf galaxy is predicted to expel $5$ to $10$ times its residual mass in stars. Dwarf galaxies can therefore pollute the IGM with metals to a much higher level than might be inferred from their stellar content. The high metallicity of gas in the central regions of clusters $\\sim Z_\\odot/3$ ({\\it e.g.} Mushotsky and Loewenstein 1997) may require a top-heavy IMF and gas ejection from massive galaxies.   \\vskip 0.2 truein  \\noindent {\\bf Aknowledgments:} The author aknowledges the award of a  PPARC Senior Fellowship."
    },
    "0002/astro-ph0002073_arXiv.txt": {
        "abstract": "We present a statistical analysis of a complete sample (255)  of northern planetary nebulae (PNe). Our analysis is based on morphology as a main parameter. The major morphological classes are: round (26\\% of the sample), elliptical (61 \\%), and bipolar (13 \\%) PNe. About a half of the round and 30 \\% of the elliptical PNe present multiple shells. Round PNe have higher galactic latitude ($|{\\rm b}| =$12) and galactic height ($<{\\rm z}>=$753 pc), than the elliptical ($|{\\rm b}| = {\\rm 7},~<{\\rm z}>=$308 pc) and bipolar ($|{\\rm b}| = {\\rm 3},~<{\\rm z}>=$179 pc). This possibly implies a different progenitor mass range across morphology, as a different stellar population would suggest. ",
        "introduction": "For decades now, it has been well established that PNe are the result of the evolution of low and intermediate mass stars (M$<$ 10 M$_{\\odot}$). The actual nebular formation process has been well understood since Kwok, Purton \\& Fitzgerald (1978) explained the formation of PN as the result of the interaction of a low density fast wind with a high density slow wind. The only drawback of this {\\it two wind} model is that it can not explain well the formation of asymmetric PNe. Since it has been long observed that most PNe do not have a round, regular shape, some other mechanism has to be invoked to produce asymmetry.  Mellema \\& Frank (1995) implemented an interacting wind model with an equatorial density enhancement. The fast low density wind interacts with an azimuthal dependent wind forming an asymmetrical PN. Rotation as a way to produce asymmetric PNe has been proposed by different authors (e.g. Calvet \\& Peimbert 1983; Ignace, Cassinelli \\& Bjorkman  1996;  Garc\\'{\\i}a-Segura \\etal  1999). The presence of a magnetic field is also able to convey an asymmetrical nebular evolution (e.g. Pascoli 1992; Chevalier \\& Luo 1994; Soker 1998; Garc\\'{\\i}a-Segura \\etal  1999). The common envelope evolutionary phase, typical of a close binary star, also produces the appropriate equatorial density enhancement to make the PN ejecta asymmetric; a similar effect can be obtained by the presence of a substellar object in the system (e.g. Soker 1997).  By correlating the morphological class with the different nebular and stellar parameters,  it may be possible to disclose the predominant mechanism responsible for the observed morphology.  Morphological studies of PNe have lacked a univocal classification scheme. Since the pioneer work by Curtis (1918), who discovered large structured \"haloes\" around some PNe, there have been several studies all aimed at the same goal: a better understanding on the physical meaning of nebular shapes. Greig (1971) classified PNe into 15 morphological classes, ultimately grouped in two main classes:  binebulous and circular; he found that binebulous PNe have lower galactic height distribution than circular PNe (Greig 1972). Zuckerman \\& Aller (1986) studied a sample of 108 PNe, classifying them into 16 morphological types, then regrouped the many classes into bipolar, elliptical, round, irregular, and other shapes; 50 \\% of their PN sample was bipolar, 30 \\% elliptical, 15 \\% round, and the rest irregular or other shapes. Zuckermann \\& Aller (1986) could not find any correlation between the morphological class and the C/O abundance. Balick (1987) divided the morphological classes in round, elliptical, and butterfly. He proposed an evolutionary sequence within each morphological type. Chu, Jacoby \\& Arendt (1987), studied a sample of 126 extended PNe. They found that the frequence of multiple-shell planetary nebulae (MSPN) was  50 \\%. Schwarz, Corradi \\& Stanghellini (1993) classified the Schwarz, Corradi \\& Melnick (1992) sample of southern Galactic PNe into elliptical, bipolar, pointsymmetric, irregular, and stellar shapes. On a subsample of the same catalog of PNe, Stanghellini, Corradi \\& Schwarz (1993) found that the central star distribution was different for bipolar and elliptical PNe. Corradi \\& Schwarz (1995), using a large PN sample ($\\sim$ 400 PNe), found a different galactic height distribution for elliptical (z=320 pc) and  for bipolar (z=130 pc) PNe.  Most of these classification schemes are based on incomplete or inhomogeneous samples. On this basis, Manchado \\etal  (1996) presented a complete set of northern Galactic PNe, to be analyzed for their morphological properties. In $\\S$2 we will illustrate this sample; we also discuss the completeness and the morphological classification based on the sample. $\\S$ 3 presents some of the relations found between the morphology and other nebular and stellar parameters, and includes a discussion on the possible evolutionary scheme for the different types of pointsymmetric PNe. Conclusions are in $\\S$ 4.  \\begin{figure} \\plotone{fig1d.ps}{} \\caption{From left to right, top to bottom: narrow band images of the PNe A 39, NGC 2438, IC 1295, He 2-429, He 2-437 and M 2-46 } \\label{fone} \\end{figure} ",
        "conclusions": "A proper statistical analysis of a complete sample of PNe has allowed us to classify them into elliptical, round, and bipolar, with the sub-classes of multiple-shell and pointsymmetric PNe. It was found that 60 \\% of our PN sample present an elliptical shape, while 26 \\% are round, and 13 \\% bipolar.  We use statistical distances that appear to be sound for the task. If they are indeed correct, the different scale heights that characterize each morphological class hint of different progenitor mass distribution for each class. Two evolutionary schemes are proposed for bipolar PNe and bipolar PNe with pointsymmetric structure."
    },
    "0002/hep-ph0002293_arXiv.txt": {
        "abstract": "We calculate the range of predicted values for $10$ quantities that will be measured by the Sudbury Neutrino Observatory (SNO). We use neutrino oscillation solutions (vacuum and MSW; active and sterile neutrinos) that are globally consistent with all available neutrino data and estimate realistic theoretical and experimental uncertainties.  The neutral current to charged current double ratio is predicted to be more than $9\\sigma$ from the no-oscillation solution for all of the currently favored neutrino oscillation solutions.  The best-fit oscillation solutions predict a CC day-night rate difference between $-0.1$\\% and $+12.5$\\% and a NC day-night difference $< 0.01$\\%. We present also the predicted range for the first and the second moments of the charged current electron recoil energy spectrum, the charged current, the neutral current, and the $\\nu$-$e$ scattering rates, the seasonal dependence of the charged current rate, and the double ratio of neutrino-electron scattering rate to charged current rate. ",
        "introduction": "\\label{sec:introduction}  What can one learn from measurements with the Sudbury Neutrino Observatory (SNO)~\\cite{sno}? What are the most likely quantitative results for each of the different experiments that can be carried out with SNO? The main goal of this paper is to help answer these questions by providing quantitative predictions for the most important diagnostic tests of neutrino oscillations that can be performed with SNO.  SNO is not an experiment. Like LEP and Super-Kamiokande, SNO is a series of experiments. We calculate the currently-favored range of predictions for $10$ quantities that are affected by neutrino oscillations and which SNO will measure.  For the impatient reader, we list here the quantities that are sensitive to neutrino oscillations which we investigate (definitions are given later in the text): first and second moments of the recoil energy spectrum, the charged current (CC), the neutral current (NC), and the neutrino-electron scattering rates, the difference between the day and the night rates for both the CC and the NC, the difference in the winter-summer CC rates, the neutral current (NC) to charged current (CC) double ratio, and the neutrino-electron scattering to CC double ratio.  The simultaneous analysis of all the SNO results, measured values and upper limits, will be a powerful technique for constraining neutrino oscillation parameters.  As an initial step in this direction, we analyze the combined results for five especially informative pairs of oscillation parameters.   \\subsection{SNO reactions} \\label{subsec:snoreactions}  The SNO collaboration will study  charged current (CC) neutrino absorption  by deuterium,  \\begin{equation} \\label{eq:nuabs} \\nu_e + d \\to p + p + e^-\\ , \\end{equation} neutral current (NC) neutrino disassociation of deuterium, \\begin{equation} \\label{eq:nunc} \\nu_x + d \\to n + p + \\nu'_x  \\ ,\\quad (x=e,\\mu,\\,\\tau), \\end{equation} and neutrino-electron scattering (ES),  \\begin{equation} \\label{eq:nuesc} \\nu_x + e^- \\to \\nu'_x + e^-\\ ,\\quad (x=e,\\mu,\\,\\tau) . \\end{equation} The energy of the recoil electrons can be measured for the CC reaction, Eq.~(\\ref{eq:nuabs}), and also for  the ES reaction, Eq.~(\\ref{eq:nuesc}).  For both these reactions, the operating energy threshold for the recoil electrons may be of order $5$ MeV.  The threshold for the NC reaction, Eq.~(\\ref{eq:nunc}), is $2.225$ MeV. Just as for radiochemical solar neutrino experiments, there is no energy discrimination for the NC reaction.  The Kamiokande~\\cite{kamiokande} and Super-Kamiokande experiments~\\cite{superk} have performed precision studies of solar neutrinos using the neutrino-electron scattering reaction, Eq.~(\\ref{eq:nuesc}). SNO will be the first detector to measure electron recoil energies as a result of neutrino-absorption, Eq.~(\\ref{eq:nuabs}). We have presented in Ref.~\\cite{snoshow} detailed predictions of what may be observed with SNO for the CC (absorption) reaction.   If there are no neutrino oscillations, i.e., $\\phi(\\nu_e) = \\phi({\\rm total})$, then the ratios of the event rates in the SNO detector are calculated to be approximately in the following proportions: CC:NC:ES = 2.05:1.00:0.19, i.e., the number of CC events is expected to exceed the number of $\\nu$-$e$ scattering events by about a factor of $11$. Since the NC efficiency is likely to be only about a half of either the CC or the ES efficiency~\\cite{sno} and currently favored oscillation solutions give $\\phi(\\nu_e) \\sim \\phi({\\rm total})$, the observed ratio of events in the SNO detector may actually be reasonably close to: CC:NC:ES $\\sim  2.0:0.5:0.2$.  In thinking about what SNO can do, it is useful to have in mind some estimated event rates for a year of operation.  The Super-Kamiokande event rate~\\cite{superk} for neutrino electron scattering is $0.475$ times the event rate that is predicted by the standard solar model~\\cite{bp98}. If there are no neutrino oscillations and the total solar neutrino flux arrives at earth in the form of $\\nu_e$ with a $^8$B neutrino flux of $0.475$ times the standard model flux, then one expects about $4.4 \\times 10^3$ CC events per year in SNO above a $5$ MeV threshold and about $1.1 \\times 10^3$ NC events, while there should only be about $415$ ES events. The above rates were calculated for a $5$ MeV CC and ES energy threshold and for a $50$\\% detection efficiency for NC events.  For an $8$ MeV threshold, the estimated CC rate is about $45$\\% of the rate for a $5$ MeV threshold and the ES rate is only about $28$\\% of the $5$ MeV threshold rate. For the currently favored oscillation solutions, the expected CC rates are typically of order $80$\\% of the rates cited above and the NC rates are about a factor of two or three higher.  \\subsection{What do we calculate?} \\label{subsec:whatcalculate}   In this paper, we calculate the likely range of quantities that are measurable with SNO using a representative sample of neutrino parameters from each of the six currently allowed $99$\\% C.L.  domains of two-flavor neutrino oscillation solutions.  In other words, we explore what can be learned with SNO, assuming the correctness of one of the six neutrino oscillation solutions~\\cite{snoshow,bks98} that is globally consistent at  $99$\\% C.L. with all of the solar neutrino experiments performed so far (chlorine~\\cite{chlorine}, Kamiokande~\\cite{kamiokande}, Super-Kamiokande~\\cite{superk}, Sage~\\cite{sage}, and GALLEX~\\cite{gallex}).  \\begin{table}[!b] \\centering \\tightenlines \\begin{minipage}{3.5in} \\caption[]{\\label{tab:bestfit} {\\bf Best-fit global oscillation parameters.} The differences of the squared masses are given in ${\\rm eV^2}$. The survival probabilities that correspond to these best-fit solutions are shown in Fig.~\\ref{fig:survival} and Fig.~\\ref{fig:survivecompare}. Results are taken from Ref.~\\cite{snoshow}} \\begin{tabular}{lcc} Scenario&$\\Delta m^2$&$\\sin^2(2\\theta)$\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} LMA&  $2.7\\times10^{-5}$ &$7.9\\times10^{-1}$\\\\ SMA& $5.0\\times10^{-6}$&$7.2\\times10^{-3}$\\\\ LOW& $1.0\\times10^{-7}$ &$9.1\\times10^{-1}$\\\\ ${\\rm VAC_S}$ & $6.5\\times10^{-11}$ &$7.2\\times10^{-1}$ \\\\ ${\\rm VAC_L}$ & $4.4\\times10^{-10}$ &$9.0\\times10^{-1}$ \\\\ ${\\rm Sterile}$ & $4.0\\times10^{-6}$ &$6.6\\times10^{-3}$\\\\ \\end{tabular} \\end{minipage} \\end{table}  Table~\\ref{tab:bestfit} lists the mixing angles and differences of mass squared for the six global best-fit solutions. Figure~\\ref{fig:survival} and Figure~\\ref{fig:survivecompare} show the survival probabilities of the best-fit solutions as a function of energy{\\footnote {For the MSW solutions, there are small but perceptible differences in the computed survival probabilities which depend upon the neutrino production probability as a function of solar radius. The survival probabilities shown in Fig.~\\ref{fig:survival} were computed by averaging the survival probability over the $^8$B production region in the BP98 model~\\cite{bp98}. In order to portray more accurately the behavior at low energy, the survival probabilities for Fig.~\\ref{fig:survivecompare} were computed by averaging over the $p-p$ production region in the BP98 model.}}   \\begin{figure}[!ht] \\centerline{\\psfig{figure=survival.ps,width=5in}} \\tightenlines \\caption[]{\\small The survival probabilities as a function of neutrino energy.  For each of the six best-fit globally acceptable neutrino oscillation solutions discussed in Ref.~\\cite{snoshow}, the figure shows the survival probability for electron type neutrinos as a function of energy. The results are averaged over one year. \\label{fig:survival}} \\end{figure}     For each measurable quantity $i$, we express our predictions based upon neutrino oscillation models in terms of the value predicted by an oscillation scenario divided by the value predicted by the combined standard electroweak model and the standard solar model. Thus for each measured quantity, $i$ (like CC or NC event rate), we evaluate the expected range of the reduced quantity $[i]$  \\begin{equation} [i] ~\\equiv~ \\frac{{\\rm (Observed~Value)_i}}{{\\rm (Standard~Model~Value)_i}}. \\label{eq:defnratio} \\end{equation} The reduced quantity $[i]$ is by construction independent of the absolute value of the solar neutrino flux, which is used in calculating both the numerator and denominator of Eq.~(\\ref{eq:defnratio}).  \\begin{figure}[!ht] \\centerline{\\psfig{figure=survivecompare.ps,width=5in,angle=270}} \\tightenlines \\caption[]{\\small Comparison of survival probabilities. For four of the best-fit oscillation solutions, the survival probabilities are compared using a linear energy scale. The differences in the energy dependence of the survival probability between the high energy region, $\\sim 8$ MeV, and the low energy region, $\\leq 1$ MeV, can be seen clearly on this figure. The MSW Sterile solution has an energy dependence similar to the SMA solution.  For clarity, we have omitted the VAC${\\rm _L}$ solution.  The parameters of the best-fit solutions are given in Table~\\ref{tab:bestfit}. \\label{fig:survivecompare}} \\end{figure}    What fluxes are used in calculating the predicted rates (e.g., for charged current or electron-neutrino interactions) implied by different neutrino scenarios? We determine the best-fit ratio of the observed neutrino flux to the standard model flux by fitting to the Super-Kamiokande rate and observed recoil electron spectrum. The procedure is described in Sec.~\\ref{sec:shapespectrum} following Eq.~(\\ref{eq:defrofe}) and Eq.~(\\ref{eq:definitionsfs}) (see especially the definition of $f(^8B)$).  We have not included explicit uncertainties in determining $f(^8B)$ for a given pair of neutrino variables, $\\Delta m^2$ and $\\sin^2 2\\theta$, but we instead have allowed $f(^8B)$ to range over the wide set of values obtained by applying our best-fit procedure at each point in the currently-allowed neutrino-parameter space.  As we shall see, the most powerful diagnostics of neutrino oscillations are formed by considering the reduced double ratio of two measurable quantities, $i$ and $j$, as follows: \\begin{equation} {[i] \\over [j]} ~\\equiv~ \\frac{{\\rm (Observed~Value)_i/(Standard~Model~Value)_i }}{{\\rm (Observed~Value)_j/(Standard~Model~Value)_j }} . \\label{eq:defndoubleratio} \\end{equation} For example, the reduced double ratio of NC to CC rates is not only independent of the absolute flux of the solar neutrinos but is also insensitive to some experimental and theoretical uncertainties that are important in interpreting the separate [NC] and [CC] rates.  We describe how we evaluate the uncertainties in Sec.~\\ref{sec:uncertainties}. All of the calculated departures from the standard model expectations are small except for the double ratio of NC to CC, [NC]/[CC]. Therefore, the theoretical and the experimental uncertainties are important.  We present in Sec.~\\ref{sec:shapespectrum} the results predicted by the six oscillation solutions for the first and second moments of the shape of the CC recoil electron energy distribution.  We summarize in Sec.~\\ref{sec:ccrate} the principal predictions for the CC rate, in Sec.~\\ref{sec:ncrate} the predictions for the neutral current rate, and in Sec.~\\ref{sec:esc} the predictions for the neutrino-electron scattering rate.  We then calculate the detailed predictions of the most important double ratios, the NC to charged current ratio, [NC]/[CC], in Sec.~\\ref{sec:ncovercc} and the neutrino-electron scattering to CC ratio, [ES]/[CC], in Sec.~\\ref{sec:esctocc}.  Up to this point in the paper, i.e., through Sec.~\\ref{sec:esctocc}, we only discuss time-averaged quantities. In Sec.~\\ref{sec:daynight}, we present the predictions for the CC of the difference between the event rate observed at night and the event rate observed during the day. For the NC rate, there is also a small difference predicted between the night rate and the day rate if the MSW Sterile solution is correct. We analyze in Sec.~\\ref{sec:seasonal} the seasonal effects in the CC rate.  Section~\\ref{sec:smokinggun} is a pairwise exploration of the discriminatory power gained by analyzing simultaneously the predictions and the observations of different smoking-gun indicators of neutrino oscillations. We consider in this section the joint analysis of variables like [NC]/[CC] versus the first moment of the CC energy spectrum, the day-night difference, or the neutrino-electron scattering rate. In Sec.~\\ref{sec:discussion}, we summarize and discuss our principal conclusions. Since we evaluate so many different effects, we give in Sec.~\\ref{sec:mostimportant} our personal list of our top four conclusions.  \\subsection{How should this paper be read?} \\label{subsec:howread}  We recommend that the reader begin by looking at the figures, which give a feeling for the variety and the size of the various quantities that can be measured with SNO. Then we suggest that the reader jump directly to the end of the paper. The main results of the paper are presented in this concluding section; the summary given in Sec.~\\ref{sec:discussion} can be used as a menu to guide the reader to the detailed analyses that are of greatest interest to him or to her.   This is the fifth in a series of papers that we have written on the potential of the Sudbury Neutrino Observatory for determining the properties of neutrino oscillations. The reader interested in details of the analysis may wish to consult these earlier works~\\cite{snoshow,howwell,bl,bkl97}, which also provide a historical perspective from which the robustness of the predictions can be judged. The present paper is distinguished from its predecessors mainly in the specificity of the predictions (representative $99$\\% C.L. predictions for each of the six currently acceptable neutrino oscillation scenarios) and in the much larger number of measurable quantities for which we now make predictions.  Recent review articles summarize clearly the present state of neutrino physics~\\cite{altarelireview,ramond} and neutrino oscillation experiments and theory~\\cite{bks98,bilenkyreview,fisher,gonzalez}.  Three and four flavor solar neutrino oscillations are discussed in Refs.~\\cite{three,four} and references cited therein.  The fundamental papers upon which all of the subsequent solar neutrino oscillation work is based are the initial study of vacuum oscillations by Gribov and Pontecorvo~\\cite{vac} and the initial studies of matter oscillations (MSW) by Mikheyev, Smirnov, and Wolfenstein~\\cite{msw}. In addition to the by-now conventional scenarios of oscillations into active neutrinos, we also consider oscillations into sterile neutrinos~\\cite{bks98,four,alexeisterile,sterilebarger,sterilevalle,sterilerabi,sterilebilenky}. ",
        "conclusions": "\\label{sec:mostimportant}  The paper contains many specific results. Here is our personal list of our most important conclusions.  (1) {\\bf The neutral current to charged current double ratio.} All currently favored active neutrino oscillation solutions predict a value for the double ratio, [NC]/[CC], of neutral current to charged current event rates that is, with our best estimates for the theoretical and experimental uncertainties, more than $9\\sigma$ away from the no-oscillation solution (neglecting statistical uncertainties).  If statistical uncertainties are included for $5000$ CC events (and $1219$ NC events), then the minimum discrepancy is reduced to $6\\sigma$.  (2) {\\bf Day-Night differences in the CC rate.} Large differences are predicted between the day and the night CC rates for some currently favored MSW solutions.  For a $5$ MeV electron recoil energy threshold, the best-fit differences vary between $-0.1$\\% (MSW Sterile) and $12.5$\\% (LMA). The largest predicted value among all the currently allowed solutions is $28.5$\\% (LMA), which could be detectable in the first year of operation of SNO. Similar results are predicted for an $8$ MeV recoil energy threshold. Vacuum oscillations have average day-night differences of order $1$\\% or even less. Small values ($< 1$\\%) of the CC day-night rate difference would be consistent with any of the six currently favored two-flavor oscillation solutions.  The day-night difference of the NC is predicted to be $< 0.01$\\% for all solutions (and is non-zero for the MSW Sterile solution).  (3) {\\bf Uncertainties.} The uncertainties in the absolute values of the neutrino cross sections for the CC and for the NC current reactions are the largest known uncertainties.  These uncertainties limit the interpretation of the separate CC and NC rates, but cancel out (to an accuracy of better than $1$\\%) of the [NC]/[CC] ratio.   (4) {\\bf  Specrum distortion.} The first moment of the CC electron recoil energy spectrum describes well the predicted deviations for all except the ${\\rm VAC_L}$ solution. For all the MSW solutions and for the ${\\rm VAC_S}$ solution, the predicted spectrum distortion is smooth and monotonic in the region accessible to SNO (see Fig.~\\ref{fig:sixspectra}) and hence can be characterized by a single parameter. The currently favored oscillation solutions predict a range of deviations of the first moment, most of which are less than the estimated $3\\sigma$ experimental uncertainty. The largest predicted deviations are for the ${\\rm VAC_S}$ (best-fit predicted deviation $283$ keV) and SMA (best-fit predicted deviation $218$ keV) solutions.  The ${\\rm VAC_L}$ solution generically predicts a bump and a dip in the low and middle energy region of the SNO electron recoil energy spectrum (see Fig.~\\ref{fig:sixspectra}). For the best-fit ${\\rm VAC_L}$ solution, this modulation is about $30$\\% and if observed would be strong evidence for the ${\\rm VAC_L}$ scenario. The predicted modulation occurs in a region where the event rate is expected to be relatively high.  (5) {\\bf Characteristic size of effects.} The current best-fit global neutrino oscillation solutions typically predict small effects, of order several percent or less, for all of the quantities that are sensitive to oscillations which SNO will measure, except [NC]/[CC].  However, for some allowed oscillation solutions, the difference between the day and the night rates and the distortion of the shape of the CC electron recoil energy spectrum may be relatively large.    (6) {\\bf Sterile neutrinos.} The current best-fit MSW Sterile solution predicts, relative to the no-oscillation solution, a $1.8$\\% shift in the first moment of the CC electron recoil energy spectrum and a CC rate that is larger than for the other currently allowed oscillation solutions. The neutral current rate is predicted to be $0.465 \\pm 0.01$ of the standard solar model rate, i.e., the MSW Sterile solution predicts a much smaller value for the neutral current rate than the other allowed oscillation solutions. The sterile solution also predicts a small but non-zero value for the difference between the NC rate during the day and the NC rate at night. The CC day-night difference is predicted to be small, but not as small as for the NC day-night difference. It will be difficult to discriminate with SNO between the no-oscillation solution and the currently allowed MSW Sterile solution.   The Sudbury Neutrino Observatory will enrich particle physics with measurements of many effects that are sensitive, in different ways, to neutrino oscillations .  We are grateful to colleagues in the SNO collaboration who have by their important experimental work and by their stimulating comments raised the questions that this paper addresses.  We are indebted to E. Akhmedov, E. Beier, S. Bilenky, D. Cowan, E. Kearns, J. Feng, M. Fukugita, K. Kubodera, E. Lisi, A. McDonald, and Y. Nir for valuable comments on a draft copy of this manuscript. JNB and AYS acknowledge partial support from NSF grant No. PHY95-13835 to the Institute for Advanced Study and PIK acknowledges support from NSF grant No. PHY95-13835 and NSF grant No. PHY-9605140."
    },
    "0002/astro-ph0002028_arXiv.txt": {
        "abstract": "{\\it Chandra} X-ray Observatory, (formerly known as AXAF), will observe down to the flux limit of 2$\\times 10^{-16}$ erg~s$^{-1}$~cm$^{-2}$. In its first year of operation {\\it Chandra}'s CCD detectors will observe over 1500 quasars serendipitously in the soft (0.5--3.5 keV) band. Over 200 quasars will be detected in X-rays in the redshift range $3<z<4$ and over 400 quasars in $2<z<3$. This will enable us to determine the high redshift X-ray luminosity function. This is the contribution by unabsorbed sources only. The total numbers would be larger by $\\sim 60$\\%. ",
        "introduction": "NASA's {\\it Chandra} X-ray Observatory was launched on July 23, 1999. The {\\it Chandra} Multiwavelength Project (ChaMP) will combine radio to X-ray observations of serendipitous {\\it Chandra} sources, with emphasis on optical identification. The ChaMP is superior to previous X-ray surveys because of (1) unprecedented X-ray positional accuracy ($\\sim 1^{\\prime\\prime}$), (2) X-ray flux limits 20 times deeper than current wide area surveys (down to $f(0.5-3.5 keV)\\sim 2\\times 10^{-16}$ erg~s$^{-1}$~cm$^{-2}$), (3) larger sky coverage ($\\sim$ 8 deg$^2$) per year than current deep surveys. ",
        "conclusions": ""
    },
    "0002/astro-ph0002502_arXiv.txt": {
        "abstract": "With a view of applications to the simulations of supernova explosion and proto neutron star cooling, we derive the Boltzmann equations for the neutrino transport with the flavor mixing based on the real time formalism of the nonequilibrium field theory and the gradient expansion of the Green function. The relativistic kinematics is properly taken into account. The advection terms are derived in the mean field approximation for the neutrino self-energy whiles the collision terms are obtained in the Born approximation. The resulting equations take the familiar form of the Boltzmann equation with corrections due to the mixing both in the advection part and in the collision part. These corrections are essentially the same as those derived by Sirera et al. for the advection terms and those by Raffelt et al. for the collision terms, respectively, though the formalism employed here is different from theirs. The derived equations will be easily implemented in numerical codes employed in the simulations of supernova explosions and proto neutron star cooling. ",
        "introduction": "The neutrino transport plays an important role in some astrophysical phenomena such as supernova explosions and the following proto neutron star cooling (e.g. \\cite{bc89,gr96,br85,su94} and references therein). In their studies, the Boltzmann equation or its approximate versions are commonly employed to describe the temporal variations of neutrino distributions in the phase space. These equations are usually derived from the following assumptions\\cite{lq66,mm89,sy99}: (1) the neutrinos are propagating along the geodesics for a massless particle, $p^{\\mu}p_{\\mu} = 0$, and the volume in the phase space occupied by these neutrinos is not varied along their world line if there is no reaction. (2) the variation of the neutrino population due to reactions is described by the so-called collision terms obtained with the stohszahl ansatz. With the masses non-diagonal in the neutrino flavor space, the neutrino oscillation occurs among different flavors of neutrinos(e.g. \\cite{fy94} and references therein). It is thus interesting from an academic point of view how this oscillation phenomenon is described by the generalized Boltzmann equations\\cite{ru90,rs93a,rs93b,pa95,dn96,sp99}. It is also important from a practical point of view for those who are interested in the possible significant consequences the oscillation might give in astrophysical events\\cite{lo95,be97,ks96,jr99}. In collapse-driven supernova explosions, for example, this is particularly the case if the resonance of oscillation occurs near a neutrino sphere where neutrinos are interacting with other particles and thus the oscillation should be treated simultaneously with these reactions and possibly with the evolution of the matter distribution as well. The purpose of this paper is to provide the formulation which can be easily implemented in those numerical simulations. \\par In considering the transport equation with the oscillation, we have to rely on a more formal derivation of the Boltzmann equation. This might be done in a couple of ways. Sirera and P\\'{e}rez\\cite{sp99}, for instance, based their derivation on the relativistic Wigner function approach in the mean field approximation. Although they took the relativistic kinematics properly into account, they did not obtain the collision terms, since it is difficult to go beyond the mean field approximation in their formalism. Raffelt et al.\\cite{rs93a,rs93b}, on the other hand, obtained their transport equation via the density matrix approach. Although they derived the collision terms, they did not consider the spatially inhomogeneous system. In this paper, we derive the relativistic Boltzmann equation including corrections due to the oscillation both in the advection terms and the collision terms by employing the real time formalism of the nonequilibrium field theory\\cite{sw61,kd65,ch85}. In this approach, the dispersion relation and the collision terms are derived on the same basis, that is, a particular approximation for the self-energy of neutrinos, which is conveniently represented with Feynman diagrams. \\par The paper is organized as follows. We first derive a generic form of the transport equation without specifying particular equations of motion of fields. Then, the formulation is applied to the neutrino flavor oscillations. In so doing, we ignore small corrections of the order of $m^{2}_{\\nu}/E_{\\nu}$ except for the terms responsible for the flavor conversion, as is usually the case. Here $m_{\\nu}$ and $E_{\\nu}$ are typical mass and energy of neutrinos in the observer's inertial frame. In this limit, as shown later, the left handed neutrinos are decoupled from the right handed ones and the difference between Majorana mass and Dirac mass never shows up in the flavor mixing. The general relativistic corrections are obtained up to the leading order of $\\lambda_{\\nu} / R$, where $\\lambda_{\\nu}$ is a typical wave length of neutrino and $R$ is a scale height of the background matter distribution. ",
        "conclusions": "With a view of application to the simulations of supernova explosion and proto neutron star cooling, we have derived a Boltzmann equation with the neutrino flavor mixing being taken into account. The derivation is based on the nonequilibrium field theory, and the ordinary gradient expansion has been performed. We assumed that the typical neutrino wave length is much shorter than the scale height of the background matter distribution, which is true for the supernova cores and proto neutron stars. The neutrino distribution matrix which is non-diagonal in the neutrino flavor space is introduced. Following the common practice, the advection part has been obtained in the mean field approximation, where the self-energy of neutrino is non-diagonal in the flavor space. This self-energy gives rise to the term in the advection part, which is responsible for the neutrino mixing and does not appear in the ordinary transport equation. The collision terms, on the other hand,  have been calculated in the Born approximation. The collision terms also have corrections due to the mixing. In these derivations, the relativistic kinematics is taken into consideration. We have further simplified the Boltzmann equation for the adiabatic flavor mixing, which is a good approximation in the supernova cores and proto neutron stars. The advection terms thus derived are essentially the same as those derived by Sirera and P\\'{e}rez\\cite{sp99}, although they employed the Wigner function formalism in the mean field approximation and did not give collision terms. The collision terms derived here, on the other hand, have the same structure as those found by Raffelt et al.\\cite{rs93a,rs93b} in the non-relativistic density matrix method. We have also shown the general relativistic correction term which accounts for the red shift and ray bending in the gravitational field and is commonly taken into account in the supernova and proto neutron star simulations.  The applications of the Boltzmann equation found here remain to be done. Since the corrections due to the flavor mixing are rather minor, particularly in the case of the adiabatic mixing, it will be simple to implement them in the neutrino transport code we have now at our disposal\\cite{sy99}. This is already underway. Since the mixing angle in matter is dependent on the neutrino energy and the direction of momentum with respect to the magnetic field if it exists. In the analyses of the neutrino flavor mixing in the supernova core, it is usually assumed that the neutrinos are flowing out radially\\cite{ks96,jr99}. However, they have an angular distribution near the neutrino sphere. Different positions of the resonant conversion due to different directions of flight of neutrinos will lead to the reduction of the neutrino flavor conversion. This will also be true in the absence of the magnetic field if the energy distribution of neutrinos and the coupling between neutrinos with different energies are taken into account. These possibilities and their implications to the mechanism of the supernova explosion, kick velocity of pulsars, and nucleosynthesis of heavy elements will be studied in the forthcoming papers."
    },
    "0002/astro-ph0002444_arXiv.txt": {
        "abstract": "We use high-resolution UV spectra taken by the Space Telescope Imaging Spectrograph instrument on board the {\\em Hubble Space Telescope} to study the 5.5 pc line of sight to the K0 V star 36~Oph~A.  The one interstellar component detected for this line of sight has a velocity inconsistent with the local interstellar cloud (LIC) flow vector, but consistent with the flow vector of the so-called G cloud, a very nearby warm cloud in the Galactic Center direction.  From our data, we measure the following values for the interstellar temperature, nonthermal velocity, H~I column density, and D/H value:  $T=5900\\pm 500$ K, $\\xi=2.2\\pm 0.2$ km~s$^{-1}$, $\\log N_{\\rm H}=17.85\\pm 0.15$, and ${\\rm D/H}=(1.5\\pm 0.5)\\times 10^{-5}$. The temperature of the G cloud is somewhat lower than that of the LIC, and Mg and Fe depletions are also lower, but the D/H value appears to be the same. Based on upper limits for the LIC absorption, we estimate the distance to the edge of the LIC to be $d_{edge}<0.19$ pc, which the Sun will reach in $t_{edge}<7400$ yrs.  The H~I Lyman-$\\alpha$ absorption line has properties inconsistent with those of the other absorption lines, indicating the presence of one or more absorption components not seen in the other lines.  We present evidence that hot neutral hydrogen local to both the Sun and the star is responsible for the excess Lyman-$\\alpha$ absorption.  This hot H~I is created by the interaction between the ISM and the winds of the Sun and 36~Oph~A.  The observed line of sight lies only $12^{\\circ}$ from the upwind direction of the LIC flow vector, where hydrodynamic models of the heliosphere suggest that heliospheric H~I absorption should be particularly prominent.  The properties of the heliospheric absorption ($T=38,000\\pm 8000$ K and $\\log N_{\\rm H}=14.6\\pm 0.3$) are consistent with previous measurements of this absorption for the $\\alpha$~Cen line of sight $52^{\\circ}$ from the upwind direction. ",
        "introduction": "In 1996, high quality spectra of the very nearby star $\\alpha$~Cen taken by the Goddard High Resolution Spectrograph instrument on board the {\\em Hubble Space Telescope} (HST) yielded a serendipitous detection of neutral hydrogen in the outer heliosphere \\citep[][hereafter LW96]{jll96}, thereby providing a unique new way to observationally study the structure and internal properties of the heliosphere.  Previous studies of heliospheric neutral hydrogen had relied primarily on Lyman-$\\alpha$ backscatter measurements, which are mostly sensitive to very nearby H~I \\citep{jlb85,eq95}. In the $\\alpha$~Cen data, the heliospheric H~I produces an absorption signature in the stellar Lyman-$\\alpha$ line that allows it to be detected despite being highly blended with absorption from the local interstellar medium (LISM).  LW96 demonstrated that the properties of the heliospheric H~I inferred from the data are consistent with the predictions of heliospheric models \\citep{vbb93,vbb95,hlp95,gpz96,gpz99}.  \\citet{kgg97} confirmed this result by directly comparing the data with H~I absorption predicted by the models, and they also found evidence that H~I in the ``astrosphere'' of $\\alpha$~Cen is also accounting for some of the non-LISM absorption observed in the Lyman-$\\alpha$ line.  The $\\alpha$~Cen line of sight lies $52^{\\circ}$ from the upwind direction of the interstellar flow through the heliosphere.  In the upwind side, the heliospheric H~I absorption is expected to be dominated by neutral hydrogen in the ``hydrogen wall'' (or ``H-wall'' for short), located between the heliopause and solar bow shock.  In this region, interactions with solar wind protons heat and decelerate the plasma component of the interstellar wind.  Charge exchange with this plasma then heats, compresses, and decelerates the neutral hydrogen as well.  The high temperature and decelerated velocity of the H~I in the hydrogen wall both play important roles in allowing absorption from this material to be detectable despite being blended with an LISM absorption component with a much larger column density.  The high temperature broadens the H-wall absorption substantially, and the deceleration shifts the absorption away from the LISM absorption, resulting in a noticeable excess of Lyman-$\\alpha$ absorption on the red side of the line.  The deceleration should be at its largest in the upwind direction, so the Lyman-$\\alpha$ signature of the H-wall should become even easier to detect if one looks closer to the upwind direction than the $\\alpha$~Cen line of sight $52^{\\circ}$ away.  Thus, we searched for another target for HST that would be much closer to the upwind direction.  We chose the nearby ($d=5.5$ pc) K0 V star 36~Oph~A (=HD 155886), which is only $12^{\\circ}$ from the upwind direction.  Our main goals in observing this star are to confirm the detection of the solar hydrogen wall claimed in the analysis of the $\\alpha$~Cen data, and to provide additional constraints on the properties of the H-wall. ",
        "conclusions": "We have analyzed absorption features observed in HST/STIS spectra of the 5.5 pc line of sight to 36~Oph~A.  Our findings are summarized as follows: \\begin{description} \\item[1.] Only one LISM absorption component is seen toward 36~Oph, with a velocity consistent with the flow vector of the G cloud.  No absorption whatsoever is detected from the LIC, meaning the edge of the LIC must be very close.  Because the line of sight is very close to the upwind directions of both the G and LIC clouds, the velocity of $-28.4$ km~s$^{-1}$ that we observe is very close to the actual flow speed of the observed material.  The fact that this {\\em projected} velocity is greater than the 26 km~s$^{-1}$ velocity of the LIC provides particularly strong evidence for the existence of faster G cloud material in that direction. \\item[2.] We estimate that the LIC D~I column density must be less than $9\\times 10^{11}$ cm$^{-2}$ to explain why no LIC absorption is detected. This corresponds to $N_{\\rm H~I}<6\\times 10^{16}$ cm$^{-2}$ based on the known LIC D/H ratio ($1.5\\times 10^{-5}$).  Assuming a density of $n_{\\rm H~I}=0.1$ cm$^{-3}$, the distance to the edge of the LIC is $d_{edge}<0.19$ pc.  The Sun will reach the edge in $t_{edge}<7400$ yrs, based on the LIC velocity vector. \\item[3.] The temperature ($T=5900\\pm 500$ K) and relative abundances of D~I, Mg~II, and Fe~II (e.g., ${\\rm D~I/Mg~II}=0.9\\pm 0.4$) are within the error bars of the measurements made for the $\\alpha$~Cen line of sight, which also samples G cloud material (LW96).  These measurements suggest that the G cloud has a slightly lower temperature and less Mg and Fe depletion than the LIC.  However, the nonthermal velocity toward 36~Oph, $\\xi=2.2\\pm 0.2$ km~s$^{-1}$, is higher than the $\\alpha$~Cen measurement ($\\xi=1.25\\pm 0.25$ km~s$^{-1}$). \\item[4.] Using a Lyman-$\\alpha$ bisector technique first described by \\citet*{bew96}, we estimate a column density of $\\log N_{\\rm H}=17.85\\pm 0.15$ toward 36~Oph and a deuterium-to-hydrogen ratio for the G cloud of ${\\rm D/H}=(1.5\\pm 0.5)\\times 10^{-5}$. The latter quantity is consistent with the LIC value of ${\\rm D/H}=(1.5\\pm 0.1)\\times 10^{-5}$. \\item[5.] The average G cloud density toward 36~Oph is $n_{\\rm H}=0.03-0.06$, significantly less than the $n_{\\rm H}\\approx 0.1$ cm$^{-3}$ densities typically observed toward the nearest stars, including $\\alpha$~Cen within the G cloud.  The most likely interpretation is that the G cloud H~I density decreases toward 36~Oph. \\item[6.] The H~I Lyman-$\\alpha$ absorption line is redshifted and suggests a much higher temperature than the other lines.  LW96 used a very similar discrepancy toward $\\alpha$~Cen to argue for the presence of absorption from heated, decelerated material in the solar hydrogen wall.  Our results strongly support this interpretation. \\item[7.] In two component fits to the H~I Lyman-$\\alpha$ line, the solar hydrogen wall component is not decelerated relative to the projected LIC velocity, which is a serious contradiction with theoretical expectations, so we propose that absorption from astrospheric material around 36~Oph is responsible for some of the absorption on the blue side of the line.  Thus, our best model of the Lyman-$\\alpha$ absorption has three components:  LISM (G cloud), heliospheric, and astrospheric. In the three component fits, the temperature and column density of the heliospheric absorption ($T=38,000\\pm 8000$ K and $\\log N_{\\rm H}=14.6\\pm 0.3$) are consistent with the measurements made for the $\\alpha$~Cen line of sight. \\end{description}"
    },
    "0002/astro-ph0002391_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "Analysis of the galactic rotation curves show that the missing Galactic Dark Matter can be explained in terms of a Cosmological Constant.\\\\ \\\\ This contribution can be considered a prime candidate for the \"Dark Energy\" which is smoothly distributed throughout space, and contributes approximately $70\\%$ to the mass/energy of the Universe \\cite{turner}.\\\\ \\\\ However, in order to support this thesis for the Cosmological Constant, thorough and systematic analysis of galactic velocity rotation data needs to take place.\\\\ \\\\ It was shown how, within the Weak Field Approximation, that  VRC data inevitably lead to a negative value for the Cosmological Constant in direct disagreement with the type Ia Supernovae data. Nevertheless, given the uncertainties in determining the deceleration parameter $q_0$ \\cite{HOYLE}, from the redshift-magnitude Hubble diagram using Type Ia supernovae as standard candles, we believe our approach is worth further consideration.\\\\ \\\\ The experimental values determined for the Cosmological Constant are shown to lie within an acceptable range. These values, used within the Extended Large Number Hypothesis, predict  values for the Quark-Hadron phase transition energy in the range 295-410 MeV.\\\\ \\\\ It would be remarkable, if proved correct, that the Cosmological Constant could be directly determined from the analysis of galactic velocity rotation curves.\\\\ \\\\ Equally remarkable, if proved correct, is the idea that astronomical observations can shed light on the last quantum mechanical phase transition of the Universe, namely the Quark-Hadron.\\\\ \\\\"
    },
    "0002/astro-ph0002114_arXiv.txt": {
        "abstract": "We present a simple spin evolution model that predicts that rapidly rotating accreting neutron stars will mainly be confined to a narrow range of spin-frequencies; $P= 1.5-5$~ms. This is in agreement with current observations of both neutron stars in the Low-Mass X-ray Binaries and millisecond radio pulsars. The main ingredients in the model are: i) the instability of r-modes above a critical spin rate, ii) thermal runaway due to heat released as viscous damping mechanisms counteract the r-mode growth, and iii) a revised estimate of the strength of dissipation due to the presence of a viscous boundary layer at the base of the crust in an old and relatively cold neutron star.  We discuss the gravitational waves that are radiated during the brief r-mode driven spin-down phase.  We also briefly touch on how the new estimates affects the predicted initial spin periods of hot young neutron stars. ",
        "introduction": "The launch of the Rossi X-ray Timing Explorer (RXTE) in 1995 heralded a new era in our understanding of neutron star physics.  Detailed observations of quasiperiodic phenomena at kHz frequencies in more than a dozen Low-Mass X-ray Binaries (LMXB) strongly suggest that these systems contain rapidly spinning neutron stars (for a recent review, see \\citet{vanderklis00}), providing support for the standard model for the formation of millisecond pulsars (MSP) via spin-up due to accretion.  Despite these advances several difficult questions remain to be answered by further observations and/or theoretical modeling. For example, we still do not know the reason for the apparent lack of radio pulsars at shorter periods than the 1.56~ms of PSR1937+21 (for a review of recent progress in the modelling of rotating neutron stars, see \\citet{stergioulas}).  The recent RXTE observations provide a further challenge for theorists. Various models suggest that the neutron stars in LMXB spin rapidly, perhaps in the narrow range 260-590~Hz \\citep{vanderklis00}.  Three different models have been proposed to explain this surprising result.  The first model (due to \\citet{white97}) is based on the standard magnetosphere model for accretion induced spin-up, while the remaining two models are rather different, both being based on the idea that gravitational radiation balances the accretion torque. In the first such model for the LMXB (proposed by \\citet{bildsten98} and recently refined by \\citet{ushomirsky00}), the gravitational waves are due to a quadrupole deformation induced in the deep neutron star crust because of accretion generated temperature gradients. The second gravitational-wave model relies on the recently discovered r-mode instability (see \\citet{akreview} for a review) to dissipate the accreted angular momentum from the neutron star.  In this Letter we reexamine the idea that gravitational waves from unstable r-modes provide the agent that balances the accretion torque. This possibility was first analyzed in detail by \\citet{akst99} (but see also \\citet{bildsten98}). Originally, it was thought that an accreting star in which the r-modes were excited to a significant level would reach a spin-equilibrium, very much in the vein of suggestions by \\citet{papaloizou} and \\citet{wagoner}.  Should this happen, the neutron stars in LMXB would be prime sources for detectable gravitational waves.  However, as was pointed out by \\citet{levin99} and \\citet{spruit}, the original idea is not viable since, in addition to generating gravitational waves that dissipate angular momentum from the system, the r-modes will heat the star up (via the shear viscosity that counteracts the r-mode at the relevant temperatures).  Since the shear viscosity gets weaker as the temperature increases, the mode-heating triggers a thermal runaway and in a few months the r-mode would spin an accreting neutron star down to a rather low rotation rate. Essentially, this conclusion rules out the r-modes in galactic LMXB as a source of detectable gravitational waves, since they will only radiate for a tiny fraction of the systems lifetime.  Other recent results would (at first sight) seem to emphasize the conclusion that the r-modes are not relevant for the LMXB. \\citet{bildsten99} investigated the effect that the presence of a solid crust would have on the r-mode oscillations.  They estimated that the dissipation associated with a viscous boundary layer that arises at the base of the solid crust in a relatively cold neutron star would greatly exceed that of the standard shear viscosity. Thus, Bildsten and Ushomirsky concluded that the r-mode instability would only be relevant for very high rotation rates, and could therefore not play a role in the LMXB.  We have reassessed the effect of the viscous boundary layer (correcting an erring factor in the estimates of \\citet{bildsten99}). Our new estimates show that the presence of the crust is important, but that the instability operates at significantly lower spin rates than suggested by Bildsten and Ushomirsky.  Once we combine our estimates with the thermal runaway (now due to heating caused mainly by the presence of the viscous boundary layer), that results as the star is spun up to the point at which the instability sets in, we arrive at a model for the spin-evolution of accreting neutron stars. Remarkably, this simple model agrees well with existing observations of rapidly rotating neutron stars, covering both the LMXB and MSP populations. ",
        "conclusions": ""
    },
    "0002/astro-ph0002322_arXiv.txt": {
        "abstract": "The powerful combination of the {\\it Chandra} X-ray telescope, the SCUBA submillimetre-wave camera and the gravitational lensing effect of the massive galaxy clusters A\\,2390 and A\\,1835 has been used to place stringent X-ray flux limits on six faint submillimetre SCUBA sources and deep submillimetre limits on three {\\it Chandra} sources which lie in fields common to both instruments. One further source is marginally detected in both the X-ray and submillimetre bands. For all the SCUBA sources our results are consistent with starburst-dominated emission. For two objects, including SMMJ\\,14011+0252 at $z=2.55$, the constraints are strong enough that they can only host powerful active galactic nuclei if they are both Compton-thick and any scattered X-ray flux is weak or itself absorbed. The lensing amplification for the sources is in the range 1.5--7, assuming that they lie at $z\\gs 1$. The brightest detected X-ray source has a faint extended optical counterpart ($I\\approx22$) with colours consistent with a galaxy at $z\\simeq 1$. The X-ray spectrum of this galaxy is hard, implying strong intrinsic absorption with a column density of about $10^{23}\\psqcm$ and an intrinsic (unabsorbed) 2--10~keV luminosity of $3\\times 10^{44}\\ergps$. This source is therefore a Type-II quasar. The weakest detected X-ray sources are not detected in {\\it HST} images down to $I\\simeq 26$. ",
        "introduction": "The spectrum of the X-ray Background (XRB) over the 1-7~keV band is a power-law of energy index 0.4 (Gendreau et al.\\ 1995) which is flatter than that of any known class of extragalactic source. Following the original suggestion of Setti \\& Woltjer (1989), it is commonly assumed that this is due to the XRB being dominated by many absorbed sources. These sources, with different absorbing column densities and redshifts, combine to give the observed XRB spectrum (Madau, Ghisellini \\& Fabian 1994; Matt \\& Fabian 1994; Comastri et al.\\ 1995; Wilman \\& Fabian 1999). The absorbed energy is presumably reradiated in the Far InfraRed (FIR) band, as observed in nearby heavily absorbed Active Galactic Nuclei (AGN) like that in the luminous {\\it IRAS} galaxy NGC\\,6240 (Vignati et al.\\ 1999). The residual harder X-ray emission which penetrates the absorbing medium, and the tail of the reradiated emission in the submillimetre band from such an source both exhibit a negative K-correction. This means that obscured AGN should be detectable to large redshifts in both bands.  Observations of the XRB with {\\it Chandra} show that much (more than 80 per cent) of it is resolved into point sources in the 2--8~keV band (Brandt et al.\\ 2000; Mushotzky et al.\\ 2000). More than half the intensity is due to a combination of hard sources identified with either otherwise normal bright galaxies or in optically faint or even invisible galaxies.  Deep 850-$\\mu$m submillimetre observations with SCUBA show that much (more than 80 per cent) of the submillimetre background seen by {\\it COBE}-FIRAS (Fixsen et al.\\ 1998) at this wavelength is resolved into discrete sources (Blain et al.\\ 1999).  A key question is what fraction of the deep X-ray and submillimetre source populations are related. The far-infrared and submillimetre background represents a significant part of the energy output of objects in the Universe. If the contribution of starbursts and AGN to the background can be separated using deep X-ray images, then the relative importance of high-mass star formation and AGN in heating the dust responsible for this emission can be determined.  Recent modelling suggests that the AGN fraction in the submillimetre background is about 20 per cent (Almaini, Lawrence \\& Boyle 1999; Fabian \\& Iwasawa 1999; Gunn \\& Shanks 1999). The uncertainty is such that AGN may contribute in total between 10--50 per cent of the total energy output of stars (Fabian 1999).  Smail et al.\\ (1997, 1998) have used the SCUBA instrument at the JCMT (Holland et al 1999) to study the submillimetre sources in seven massive clusters of galaxies at redshifts between 0.2 and 0.4, exploiting the gravitational lensing magnification from the clusters to make an exceptionally deep 850-$\\mu$m survey for background sources. {\\it Chandra} images of two of the clusters, A\\,2390 and A\\,1835, have recently been obtained which, owing to the superb angular resolution, probe much deeper than any previous X-ray images of these regions, for example the {\\it ROSAT}-HRI limit of $8 \\times 10^{-14}$\\,erg\\s$^{-1}$\\,cm$^{-2}$ to the 0.1--2.0\\,keV X-ray flux of the SCUBA source SMMJ\\,14011+0252 in A\\,1835 (Ivison et al.\\ 2000). Here we study the {\\it Chandra} X-ray limits to the flux of the SCUBA sources and conversely SCUBA limits on the {\\it Chandra} sources found in the images. Only one source is marginally detected in both wavebands. We then discuss the implications for obscured AGN models. ",
        "conclusions": "We compare expected values of $\\alpha_{\\rm SX}$ as a function of redshift in Fig.~2 for various classes of observed objects, both starbursts and AGN. We select 3C\\,273 as an example of a powerful quasar, and use the submillimetre data of Neugebauer, Soifer \\& Miley (1985) in our model; Arp\\,220 as a starburst using the submillimetre data of Sanders et al.\\ (1999) and the X-ray data of Iwasawa (1999); and NGC\\,6240 as a powerful obscured (Compton-thick) AGN using an estimate of the submillimetre spectrum based on that from Arp\\,220 normalised by the FIR luminosity and the {\\it BeppoSAX} spectrum from Vignati et al.\\ (1999). We also show how $\\alpha_{\\rm SX}$ changes as the absorption in NGC\\,6240 decreases from the measured value of $2\\times 10^{24}$ to $5\\times 10^{23}\\psqcm$ and also if the scattered fraction drops to 1 per cent.  Data from several quasars at $z>4$ are shown for comparison with the 3C\\,273 prediction.  \\begin{figure} \\centerline{\\psfig{figure=fig2.ps,width=0.5\\textwidth,angle=270}} \\caption{Submillimetre to X-ray spectral index, $\\alpha_{\\rm SX}$, plotted against redshift for a powerful quasar, 3C\\,273, an absorbed AGN, NGC\\,6240, and a starburst, Arp\\,220. The effects of decreasing the scattered fraction from NGC\\,6240 to 1 per cent, and of its column density to $5\\times 10^{23}\\psqcm$ are indicated by the dash-dot and dash lines. The {\\it Chandra}/SCUBA sources with known redshifts are indicated by the squares. The $\\ast$ symbols represent the five $z>4$ quasars for which both X-ray and SCUBA (850$mu$m) data are available; monochromatic X-ray fluxes at 2 keV were calculated from the (Galactic) absorption-corrected 0.1--2.0 keV fluxes in Kaspi et al.~(2000) assuming a power-law of $\\Gamma=2$; 850-$\\mu$m fluxes are taken from the compilation in McMahon et al.~(1999).} \\end{figure}  The typical limits on $\\alpha_{\\rm SX}$ for the SCUBA galaxies are $\\alpha_{\\rm SX} \\ls 1.2$ (at 99 per cent confidence).  Combining the limits for the 6 galaxies we obtain a 99 per cent confidence limit on a typical SCUBA galaxy of $\\alpha_{\\rm SX}> 1.34$, with the strongest constraints for individual galaxies being $\\alpha_{\\rm SX} > 1.32$  for SMMJ\\,14011+0252 ($z=2.55$) and SMMJ\\,14009+0252 in the A\\,1835 field. Thus the  indices for this population are straightforwardly consistent with starbursts. They do not resemble Compton-thin AGN at any redshift and are inconsistent with the spectrum of NGC\\,6240 at any redshift. If these galaxies host a powerful AGN then either a) reprocessed radiation from the AGN provides only a minor fraction of the submillimetre luminosity, or b) the source must be Compton thick and in addition the fraction of any scattered X-ray emission must be less than one per cent.  We note that X-ray limits on the hyperluminous {\\it IRAS} galaxy F15307+3252 at $z=0.92$, which does contain an AGN since broad scattered emission lines are seen (Hines et al.\\ 1995), show that it must have a very low X-ray scattered fraction ($<1$ per cent), possibly due to intrinsic absorption of the scattered emission (Fabian et al.\\ 1996). If this is typical of obscured AGN more powerful than NGC\\,6240, then the present results would be consistent with obscured AGN provided that they are at relatively high redshift ($z\\gs 2$).  Of the four {\\it Chandra} sources we identified in our fields, two have probable optical counterparts suggesting that they are distant disk galaxies.  The colours and luminosity of the brightest of these suggest it is an $L^\\ast$ mid-type galaxy at $z\\sim 1$.  If the two optically-unidentified {\\it Chandra} sources are AGN, and the host galaxies have luminosities of $L^\\ast$ or greater, then they must be at $z>2$, or be intrinsically reddened. Obscuration of both AGN and surrounding spheroid is required by some models for the XRB (Fabian 1999).  It is possible for these dust-enshrouded strongly-obscured AGN to be undetected in our 850-$\\mu$m SCUBA maps  if their dust is typically much hotter than 40\\,K.  Comparing the radio flux and SCUBA limit for the brightest {\\it Chandra} galaxy CXOUJ215334.0+174240 (\\S2), the redshift $z$ and dust temperature $T_d$ of the source must satisfy the relationship $T_d > 27(1+z)$\\,K.  Taking the redshift constraint from the optical, $z \\simeq 0.85\\pm 0.15$, we derive $T \\ge 50$\\,K.  Such sources would not contribute substantially to the submillimetre background at $\\sim 1$\\,mm, but might at shorter wavelengths.  Eventual comparison with  {\\it ISO} imaging may shed further light on the nature of the dust emission in these sources.  In summary, {\\it Chandra} and SCUBA observations have been combined to probe the relation between the faint X-ray and submillimetre sky. Only one marginal source is seen in both datasets, and so in general we find deep submillimetre limits from SCUBA on the X-ray sources found using {\\it Chandra} and deep X-ray limits from {\\it Chandra} on the SCUBA-selected sources. The limits on background sources in these fields are particularly strong due to amplification by gravitational lensing.  For the SCUBA galaxies, we cannot completely rule out the presence of either an obscured, weak AGN or a more powerful Compton-thick AGN in which any scattered flux is also  very weak or absorbed.  However, the simplest explanation of our results on the SCUBA sources is that they are predominantly powered by starbursts.  Clearly the current sample is small and so we cannot make any definitive statement about the whole SCUBA population. We note that at the moment our conclusions are not inconsistent with suggestions that AGN powering 20 per cent of the SCUBA population, although a substantially higher fraction would be difficult to accommodate.  For the {\\it Chandra} sources in the A2390 field we find that they have optically faint counterparts, $I\\gs 23$--27. We identify one source as a probable Type-II obscured quasar at $z\\simeq 1$.  We suggest that the remaining, typically fainter sources are either more distant, $z>2$, or intrinsically reddened.  Observations of these fields in the near-infrared are urgently needed to test the nature of these sources."
    },
    "0002/astro-ph0002052_arXiv.txt": {
        "abstract": "We analyze the properties of the pulsed emission from the accreting millisecond pulsar \\saxpsr in observations of its April 1998 outburst by the Rossi X-Ray Timing Explorer. Pulse phase spectroscopy shows that the emission evolves from a hard spectrum (power law with photon index $2.39\\pm0.06$) to a soft spectrum (index $3.39\\pm0.24 $). This softening is also observable as a phase lag in the fundamental of low-energy photons with respect to high-energy photons. We show that this lag is roughly constant over ten days of the outburst.  We fit these data with a model where the pulse emission is from a hot spot on the rotating neutron star and the flux as a function of phase is calculated including the effects of general relativity.  The energy-dependent lags are very well described by this model. The harder spectra at earlier phases (as the spot approaches) are the result of larger Doppler boosting factors which are important for this fast pulsar. ",
        "introduction": "Strong X-ray pulsations with a 2.49 msec period were discovered from \\saxpsr in an April 1998 observation with the Rossi X-ray Timing Explorer \\citep{wk98}. The pulsar is a member of an accreting binary system with an orbital period of 2.01 hour and a low-mass companion \\citep{cm98}. Though similar to other low-mass X-ray binaries in its timing and spectral properties \\citep[e.g.][]{wk99,hs98}, \\saxpsr is unique for its X-ray pulsations. No other such binary has shown coherent pulsations in its persistent flux despite of careful searches \\citep[][and references therein]{vaughan94a}.  As such, \\saxpsr is the fastest rotating accreting neutron star. If the pulsations are due to modulated emission from one hot spot on the neutron star surface, the 2.49 msec period corresponds to an equatorial speed of approximately 0.1c. With these high speeds, \\saxpsr offers an excellent system for studying relativistic effects.  One such effect may be the observed lag of low-energy photons relative to high-energy photons in the pulse discovered by \\citet{cmt98}. \\citet{cmt98} suggest that the lags are due to Comptonization in a relatively cool surrounding medium. Alternatively, the lags may be the result of a relativistic effect: the high-energy photons are preferentially emitted at earlier phases due to Doppler boosting along the line of sight. This possibility was suggested for the similar lags in the 549~Hz oscillations in an X-ray burst of Aql~X-1 \\citep{ford99}, where a simple model showed that the delays roughly match those expected.  In the following, we present new measurements of the pulsed emission from \\saxpsr. We show that the energy-dependent phase lags are equivalent to a hardening pulse profile. We model this behavior in terms of a hot spot on the neutron star, including relativistic effects. ",
        "conclusions": "The pulsed emission from \\saxpsr evolves through its phase from a relatively hard to a soft spectrum, as shown by our phase resolved spectroscopy (Figure~\\ref{lc}). This evolution can also be thought of, and measured as, an energy-dependent phase lag in the fundamental (Figure~\\ref{lags}), i.e. higher-energy photons emerging earlier in phase than lower-energy photons.  We have applied a model to the data which consists of a hot spot on the rotating neutron star under a general relativistic treatment. The dominant effect accounting for energy-dependent delays is Doppler boosting, the larger boosting factors at earlier phases giving harder spectra. The model fits the data quite well.  The model also provides a stable mechanism for generating the phase delays, which meshes with the fact that the characteristic delays remain stable in time to within 25\\% (Figure~\\ref{lags_v_time}) even as there is a factor of two decrease in the X-ray flux, a possible tracer of accretion rate. As noted in \\citet{ford99}, in addition to explaining the energy-dependent phase lags in \\saxpsr, this mechanism may also account for the lags in the burst oscillations of Aql X-1 \\citep{ford99} and kilohertz QPOs \\citep{vaughan98,kaaret99a}. Phase resolved spectroscopy has not yet been possible in these signals.  The model offers a new means of measuring the neutron star mass and radius, a notoriously difficult problem in accreting binary systems. The radius is directly related to $v$, the equatorial speed ($v=\\Omega_{\\rm spin} R$). The fits also depend on the model parameters $M$ and $R/M$, though the lag spectra are less sensitive to these quantities.  The model could be independently constrained by future optical spectroscopy of the companion which could provide information on $M$ and the viewing angles, $\\alpha$ and $\\beta$."
    },
    "0002/astro-ph0002264_arXiv.txt": {
        "abstract": "I present an overview of optical observations (mostly spectra) of Type II supernovae. SNe~II are defined by the presence of hydrogen, and exhibit a very wide variety of properties. SNe~II-L tend to show evidence of late-time interaction with circumstellar material.  SNe~IIn are distinguished by relatively narrow emission lines with little or no P-Cygni absorption component and (quite often) slowly declining light curves; they probably have unusually dense circumstellar gas with which the ejecta interact. Some SNe~IIn, however, might not be genuine SNe, but rather are super-outbursts of luminous blue variables.  The progenitors of SNe~IIb contain only a low-mass skin of hydrogen; their spectra gradually evolve to resemble those of SNe~Ib. Limited spectropolarimetry thus far indicates large asymmetries in the ejecta of SNe~IIn, but much smaller ones in SNe~II-P. There is intriguing, but still inconclusive, evidence that some peculiar SNe~IIn might be associated with gamma-ray bursts. SNe~II-P are useful for cosmological distance determinations with the Expanding Photosphere Method, which is independent of the Cepheid distance scale. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/hep-ph0002090_arXiv.txt": {
        "abstract": "s{The most recent cosmological data point in the direction of a cosmological constant dominated universe.  A suitable candidate for providing the required acceleration is a rolling scalar field named ``quintessence''.  We address the issue of building a particle physics model for quintessence in the context of Susy QCD.  We then go on to ask if the quintessence scalar can be identified with the inflaton field and give two examples in which this is indeed the case.} ",
        "introduction": "The very last years have witnessed growing interest in cosmological models with $\\Omega_m \\sim 1/3$ and $\\Omega_\\Lambda \\sim 2/3$, following the most recent observational data (see for example the discussion in the paper by Bahcall {\\it et al.} \\cite{data} and references therein). A very promising candidate for a dynamical cosmological constant is a rolling scalar field, named ``quintessence''. \\cite{quint} The main motivation for constructing such dynamical schemes resides in the hope of weakening the fine tuning issue implied by the smallness of $\\Lambda$. In this respect, a very suitable class of models is provided by inverse power scalar potentials which admit attractor solutions characterized by a negative equation of state. \\cite{scalcosmo,track}  Consider the cosmological evolution of a scalar field $Q$, with potential $V(Q)=M^{4+\\alpha }Q^{-\\alpha }$, $\\alpha >0$, in a regime in which the scalar energy density is subdominant with respect to the background. Then it can be shown \\cite{scalcosmo,track} that the solution $Q\\sim t^{1-n/m}$, with $n=3(w_{Q}+1)$ and $m=3(w_{B}+1)$, is an attractor in phase space. We have defined $w_Q$ to be the equation of state of the scalar field $Q$, and $w_B$ that of the background ($=1/3$ for radiation and $=0$ for matter). The equation of state of the scalar field on the  attractor is found to be $w_{Q}= (\\alpha \\, w_B -2)/(\\alpha+2)$, which is always negative during matter domination. As a consequence, the ratio of the scalar to background energy density is not constant but scales as $\\rho _{Q}/\\rho_{B}\\sim a^{m-n}$, thus growing during the cosmological evolution, since $n$ $<m$. The behaviour of these solutions is determined by the cosmological background and for this reason they have been named ``trackers'' in the literature. \\cite{track} A good feature of these models is that for a very wide range of the initial conditions the scalar field will reach the tracking attractor before the present epoch. Depending on the initial time, you can have several tens of orders of magnitude of allowed initial values for the scalar energy density. This fact, together with the negative equation of state, makes the trackers feasible candidates for explaining the cosmological observation of a presently accelerating universe. The point at which the scalar and matter energy densities are of the same order depends on the mass scale in the potential.  This is fixed by requiring that $\\Omega_Q = {\\cal O} (1)$ today.  An interesting question, then, is whether the `quintessence' scalar and the inflaton field, which dominate the expansion of the universe at very different times, could indeed be the same field. If this is the case, it should also be possible to uniquely fix the initial conditions for the `quintessential' rolling from the end of inflation. Models in which one single scalar field drives both inflation and the late time cosmological accelerated expansion are named ``quintessential inflation'' models. \\cite{pv,noi2} ",
        "conclusions": ""
    },
    "0002/astro-ph0002270_arXiv.txt": {
        "abstract": "The formation and evolution of clusters of galaxies are sensitive to the underlying cosmological model. Constraints on cosmological parameters of cold dark matter models have been derived from mass, temperature and luminosity functions of clusters. We study the importance of including cluster formation history and a correction for collapsed fraction of objects in determining the cluster X-ray temperature function. We find that both effects are important. We compare temperature functions obtained by using a power law approximation for the mass variance normalized to X-ray clusters to those obtained by using a COBE normalized full CDM treatment. We conclude that the temperature function could be a powerful test on the average density of the Universe if we could find the correct way of interpreting the data. ",
        "introduction": "Clusters of galaxies are the largest known virialized gravitational systems in the Universe. According to our working hypothesis of structure formation, the cold dark matter (CDM) model, clusters form via gravitational collapse of high density peaks of primordial density fluctuations. The distribution and evolution of these peaks strongly depend on the CDM model parameters. We use the following parameterization for our CDM models: $\\Omega_0$, $\\Omega_b$, $\\Omega_\\Lambda$, $h$ (matter density, baryon density, cosmological constant, all in units of the critical density, all at $z$ = 0, and $h$ is the Hubble constant in units of $100 \\, \\rm Km\\,s^{-1}\\,Mpc^{-1}$). Semi-analytical methods, like the Press-Schechter approximation, use some version of a spherical collapse model to predict the distribution of collapsed objects and their masses as a function of redshift. Assuming a physical model for the clusters and their intracluster medium, one can translate cluster virial masses to X-ray luminosity, or temperature of the intracluster gas, which can be compared to observations. We use the temperature function of clusters in our study. The most popular method is to assume a power law approximation for the power spectra of CDM models, use the Press-Schechter mass function (PSMF) to predict the abundance, and then derive the virial temperature to get the intracluster gas temperature, or integrate over formation epochs to determine the mass-temperature conversion. This method has the power spectrum normalization and the power law exponent of the mass variance on cluster scales as free parameters and determines matter density, and has no dependence on the Hubble constant. We use COBE normalized spectra of CDM models to derive the mass variance and the Press-Schechter mass function to derive the distribution of cluster masses. This method returns the matter density as a function of the Hubble constant, thus one can not derive either parameter separately. We use four different methods to obtain theoretical temperature functions. We integrate over formation epochs in methods A and B but not in methods C and D. We take into account the collapsed fraction of objects in methods A and C and not in B and D. Method A thus take both effects and method D none of them into account. For each method (A, B, C, and D) we use a grid of CDM model parameters, $0.2 < \\Omega_m < 0.9$ and $0.2 < h < 0.9$, and calculate the temperature functions assuming open and flat CDM cosmologies. We compare these theoretical temperature functions to data and determine the best fit models by minimizing the corresponding $\\chi^2$. ",
        "conclusions": "In Figure~1a we show the resulting best fit temperature functions using the power law approximation of Donahue and Voit (1999) and Blanchard et al. (1999) as we reconstructed them following their method. We used Horner et al. (1999)'s compilation of data taken from Edge et al. (1990) Henry and Arnaud (1991), and Markevitch (1998) (squares, diamonds, and triangles). The solid and dashed dotted lines represent best fit flat model, $\\Omega_m = 0.27$, $\\sigma_8 = 0.73$ $\\alpha = 0.13$, of Donahue and Voit (they integrated over cluster formation epoch) and their best fit model, but not integrated over cluster formation epoch. Blanchard et al.'s best fit flat model, $\\Omega_m = 0.735$, $\\sigma_c = 0.623$, $\\alpha = 0.18$, is represented by a long dashed line. The short dashed line represents Blanchard et al.'s best fit temperature function, but integrated over cluster formation epoch. Open model temperature functions (not shown) are very similar to those of the flat models. The difference between temperature functions from methods with and without integration over cluster formation epoch is significant even with existing data. The large difference between density parameters derived by Donahue and Voit, and Blanchard et al. is due to the different normalizations of the $M(T)$ relation and not integration over cluster formation, which can cause a change less than 0.1 in the density parameter (cf. next paragraph).  Figure~1b shows temperature functions of the best fit flat CDM model (best fit using method A) with $\\Omega_m = 0.39$, $\\Lambda = 0.61$, $h = 0.5$, and this best fit model, but temperature functions with and without taking the collapsed fraction and/or integrating over cluster formation epochs (our methods B, C and D). The data are the same as in figure 1a. The solid, long dashed, short dashed and dashed dot lines represent temperature functions using our methods A, B, C and D. Open models behave similarly so we do not show them here. The smallest effect is the correction for collapsed fraction of objects which is about the size of the error bars of the data, thus an important effect (compare methods A and B). A larger difference is caused by integration over formation epoch (methods A and C). If one fits models without integrating and/or taking collapsed fraction into account (methods B, C and D), one get matter densities $\\Omega_m = 0.36$, $\\Omega_m = 0.34$, and $\\Omega_m = 0.31$ (assuming $h = 0.5$). Not taking either effects into account causes about 0.1 change in the derived matter density (compare results using methods A and D). As we can see, the temperature function is very sensitive to CDM model parameters if one uses a full CDM treatment, and thus a precise determination of the matter density is possible if the Hubble constant is known. We should keep in mind, however, that the normalization of the $M(T)$ relation, for example, causes much larger error in determining $\\Omega_m$. Since this method does not allow us to separate the density parameter from the Hubble constant, the result is a best fit function of the two. We can make use of the results from the power law approximation which gives the best fit density parameter, and check if the corresponding Hubble constant is reasonable using our methods. We find that our best fit CDM models yield the same matter density as the best fit models of Donahue and Voit (1999) and Blanchard et al. (1999) if we use $h = 0.55$ and $h = 0.35$. Thus Blanchard et al.'s method is only marginally compatible to ours.  Figure 1b shows that  systematic errors from interpretation of the temperature function are larger than the error bars on even the existing data. We conclude that, if the Hubble constant is known, a comparison between the observationally derived cluster temperature function and those derived from a full CDM treatment may yield an accurate determination of the density parameter (with an error less then 0.05), however, this can be done only if we find other ways to derive a correct theoretical model to interpret the data.    \\begin{figure} \\centerline{\\psfig{file=fig1a.eps, width=6cm} \\psfig{file=fig1b.eps, width=6cm}} \\caption[]{a) left panel: temperature functions using best fit power law approximations; b) right panel: temperature functions of the best fit CDM flat model using different methods. See text for details.} \\end{figure}"
    },
    "0002/astro-ph0002046_arXiv.txt": {
        "abstract": "We are studying star formation effects on the properties of the ISM in low metallicity environments using mid-infrared (MIR) and far-infrared (FIR) observations of starbursting dwarf galaxies taken with the Infrared Space Observatory (ISO) and the Kuiper Airborne Observatory (KAO). Effects of the hard pervasive radiation field on the gas and dust due to the dust-poor environments are apparent in both the dust and gas components. From a 158 $\\mu$m [CII] survey we find enhanced I[CII]/FIR ratios in dwarf galaxies and I[CII]/I(CO) ratios up to 10 times higher than those for normal metallicity starburst galaxies.  We consider MIR observations in understanding the star formation properties of dwarf galaxies and constraints on the stellar SED. Notably, the strong MIR [NeIII]/[NeII] ratios reveal the presence of current massive stellar populations $<$ 5 My old in NGC~1569, NGC~1140 and IIZw40. The MIR unidentified infrared bands (UIBs) are weak, if present at all, as a general characteristic in low metallicity environments, revealing the destruction of the smallest carbon particles (e.g. PAHs) over large spatial scales. This is confirmed with our dust modeling: mass fractions of PAHs are almost negligible compared to the larger silicate grains emitting in the FIR as well as the small carbon grains emitting in the MIR, which appear to be the source of the photoelectric gas heating in these galaxies, in view of the [CII] cooling. ",
        "introduction": "\\typeout{SET RUN AUTHOR to \\@runauthor} To construct a comprehensive picture of a galaxy's history, understanding the distribution of its energy budget is a fundamental step.  For this we must consider observations covering several characteristic wavelength regimes, thus, sampling the various components of the interstellar medium (ISM). While the UV to NIR wavelength continua give us relatively direct probes of the stellar populations, this radiation is subject to varying amounts of absorption before we view it. Some of this energy is absorbed by the gas directly in HII regions or transferred to the gas in photodissociation regions (PDRs) and reemitted as molecules, bands and atomic ionic and recombination lines, from wavelengths covering the UV to FIR and beyond. Some of the stellar energy is absorbed by the dust, revealed through extinction, and reradiated in MIR to submillimeter wavelengths as thermal emission.  Therefore, models of the ISM in galaxies must consider these interdependent processes and be self-consistent. Our knowledge of the wavelength window from the MIR to the FIR has been limited by the low spatial and spectral resolution provided by IRAS, and has been rather sketchy when it comes to detailed studies of the ISM of individual galaxies. The Infrared Space Observatory (ISO) \\cite{kessler96} has been a recent turning point in this effort, providing high spectral and spatial resolution and unprecedented sensitivity in the MIR to FIR.  We are incorporating our MIR and FIR observations in a study of the energy redistribution in starburst galaxies to understand the effects of the star formation on the surrounding gas and dust. Here we report the progress to date in our study of star forming low-metallicity dwarf galaxies, which, in the absence of major dynamical complications, allow us to `simplify' model assumptions and the interpretation of observations. ",
        "conclusions": "Tracers of various components of the ISM show evidence of effects of the hard stellar radiation field in dwarf galaxies on the surrounding ISM due to the decreased dust abundace, allowing photoionization over large galactic scales to occur. From our survey of the 158 $\\mu$m [CII] PDR cooling line in dwarf galaxies, we observe an increased penetration of the FUV radiation field which enhances the I[CII]/I(CO) emission in dwarf galaxies up to a factor of 10 more than in normal metallicity star burst galaxies. We also find a small enhancement in the I[CII]/FIR ratio in dwarf galaxies. Followup MIR ISOCAM spectroscopy provides details of ionic lines, UIBs and small hot small grain emission distribution in dwarf galaxies. The strong MIR [NeIII]/[NeII] ratios are signatures of the hard radiation fields and indicate the presence of young massive stellar populations in dwarf galaxies. Because of the increase in T$_{eff}$ in low metallicity environments, this ratio is enhanced at least 5 to 10 times more in dwarf galaxies than in normal metallicity galaxies. The penetrating radiation field also effects the dust components, destroying the UIBs in some dwarf galaxies on global scales, as is evident in the MIR spectra and in the dust modeling."
    },
    "0002/astro-ph0002336_arXiv.txt": {
        "abstract": "We study the expected redshift evolution of galaxy cluster abundance between $0\\lsim z \\lsim 3$ in different cosmologies, including the effects of the cosmic equation of state parameter $w\\equiv p/\\rho$.  Using the halo mass function obtained in recent large scale numerical simulations, we model the expected cluster yields in a 12~deg$^{2}$ Sunyaev-Zel'dovich Effect (SZE) survey and a deep 10$^{4}$~deg$^{2}$ X-ray survey over a wide range of cosmological parameters.  We quantify the statistical differences among cosmologies using both the total number and redshift distribution of clusters. Provided that the local cluster abundance is known to a few percent accuracy, we find only mild degeneracies between $w$ and either $\\Omega_m$ or $h$.  As a result, both surveys will provide improved constraints on $\\Omega_m$ and $w$. The $\\Omega_m$--$w$ degeneracy from both surveys is complementary to those found either in studies of CMB anisotropies or of high--redshift Supernovae (SNe).  As a result, combining these surveys together with either CMB or SNe studies can reduce the statistical uncertainty on both $w$ and $\\Omega_m$ to levels below what could be obtained by combining only the latter two data sets. Our results indicate a formal statistical uncertainty of $\\approx 3\\%$ (68$\\%$ confidence) on both $\\Omega_m$ and $w$ when the SZE survey is combined with either the CMB or SN data; the large number of clusters in the X--ray survey further suppresses the degeneracy between $w$ and both $\\Omega_m$ and $h$. Systematics and internal evolution of cluster structure at the present pose uncertainties above these levels.  We briefly discuss and quantify the relevant systematic errors.  By focusing on clusters with measured temperatures in the X--ray survey, we reduce our sensitivity to systematics such as non-standard evolution of internal cluster structure. ",
        "introduction": "It has long been realized that clusters of galaxies provide a uniquely useful probe of the fundamental cosmological parameters.  The formation of the large--scale dark matter (DM) potential wells of clusters is likely independent of complex gas dynamical processes, star formation, and feedback, and involve only gravitational physics.  As a result, the abundance of clusters $N_{\\rm tot}$ and their distribution in redshift $dN/dz$ should be determined purely by the geometry of the universe and the power spectrum of initial density fluctuations.  Exploiting this relation, the observed abundance of nearby clusters has been used to constrain the amplitude $\\sigma_{\\rm 8}$ of the power spectrum on cluster scales to an accuracy of $\\sim 25\\%$ (e.g. \\cite{white93,viana96}).  The value of $\\sigma_{\\rm 8}$ in these studies depends on the assumed underlying cosmology, especially on the density parameters $\\Omega_m$ and $\\Omega_\\Lambda$. Subsequent works (\\cite{bahcall98,blanchard98,viana99}) have shown that the redshift--evolution of the observed cluster abundance places useful constrains on these two cosmological parameters.  In the above studies, the equation of state for the $\\Lambda$--component has been implicitly assumed to be $p=w\\rho$ with $w=-1$.  The recent suggestion that $w$ might be different from $-1$, or even redshift dependent (\\cite{turner97,caldwell98}) has inspired several studies of cosmologies with a component of dark energy.  From a particle physics point of view, such $w>-1$ can arise in a number of theories (see \\cite{freese87,ratra88,turner97,caldwell98} and references therein). It is therefore of considerable interest to search for possible astrophysical signatures of the equation of state, especially those that distinguish $w=-1$ from $w>-1$. Wang et al. (2000) has summarized current astrophysical constraints that suggest $-1 \\leq w \\lsim -0.2$; while recent observations of Type Ia SNe suggest the stronger constraint $w \\lsim -0.6$ (\\cite{ptw99}).  The galaxy cluster abundance provides a natural test of models that include a dark energy component with $w\\ne -1$, because $w$ directly affects the linear growth of fluctuations $D_z$, as well as the cosmological volume element $dV/dzd\\Omega$. Furthermore, because of the dependence of the angular diameter distance $d_A$ on $w$, the experimental detection limits for individual clusters, e.g., from the Sunyaev--Zel'dovich effect (SZE) decrement or the X--ray luminosity, depend on $w$.  Wang \\& Steinhardt (1998, hereafter WS98) studied the constraints on $w$ from a combination of measurements of the cluster abundance and Cosmic Microwave Background (CMB) anisotropies.  Their work has shown that the slope of the comoving abundance $dN/dz$ between $0<z<1$ depends sensitively on $w$, an effect that can break the degeneracies between $w$ and combinations of other parameters $(h,\\Omega,n)$ in the CMB anisotropy alone.  Here we consider in greater detail the constraints on $w$, and other cosmological parameters, from cluster abundance evolution.  Our main goals are: (1) to quantify the statistical accuracy to which $w\\neq -1$ models can be distinguished from standard $\\Lambda$ Cold Dark Matter (CDM) cosmologies using cluster abundance evolution; (2) to assess these accuracies in two specific cluster surveys: a deep SZE survey (\\cite{carlstrom99}) and a large solid angle X--ray survey, and (3) to contrast constraints from cluster abundance to those from CMB anisotropy measurements and from luminosity distances to high--redshift Supernovae (\\cite{schmidt98,perlmutter99}).  Our work differs from the analysis of WS98 in several ways. We examine the surface density of clusters $dN/dzd\\Omega$, rather than the comoving number density $n(z)$.  This is important from an observational point of view, because the former, directly measurable quantity inevitably includes the additional cosmology-dependence from the volume element $dV/dzd\\Omega$.  We incorporate the cosmology--dependent mass--limits expected from both types of surveys. Because the SZE survey has a nearly $z$--independent sensitivity, we find that high--redshift clusters at $z>1$ yield useful constraints, in addition to those studied by WS98 in the range $0<z<1$. Finally, we quantify the statistical significance of differences in the models by applying a combination of a Kolmogorov--Smirnov (KS) and a Poisson test to $dN/dzd\\Omega$, and obtain constraints using a grid of models for a wide range of cosmological parameters.  This paper is organized as follows. In \\S~\\ref{sec:surveys}, we describe the main features of the proposed SZE and X-ray surveys relevant to this work. In \\S~\\ref{sec:models} we briefly summarize our modeling methods and assumptions. In \\S~\\ref{sec:sensitivity}, we quantify the effect of individual variations of $w$ and of other parameters on cluster abundance and evolution. In \\S~\\ref{sec:wconstraints}, we obtain the constraints on these parameters by considering a grid of different cosmological models. In \\S~\\ref{sec:discussion}, we discuss our results and the implications of this work.  Finally, in \\S~\\ref{sec:conclusions}, we summarize our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We studied the expected evolution of galaxy cluster abundance from $0\\lsim z \\lsim 3$ in different cosmologies, including the effects of variations in the cosmic equation of state parameter $w\\equiv p/\\rho$.  By considering a range of cosmological models, we quantified the accuracy to which $\\Omega_m$, $w$, and $h$ can be determined in the future, using a 12~deg$^{2}$ Sunyaev-Zel'dovich Effect survey and a deep 10$^{4}$~deg$^{2}$ X-ray survey.  In our analysis, we have assumed that the local cluster abundance is known accurately: we find that in practice, an accuracy of $\\sim 5\\%$ is sufficient for our results to be valid.  We find that raising $w$ significantly flattens the redshift--distribution, which can not be mimicked by variations in either $\\Omega_m$, $h$, which affect essentially only the normalization of the redshift distribution. As a result, both surveys will be able to improve present constraints on $w$.  In the $\\Omega_m-w$ plane, both the SZE and X--ray surveys yield constraints that are highly complementary to those obtained from the CMB anisotropy and high--$z$ SNe.  Note that the SZE and X--ray surveys are themselves somewhat complementary. In combination with these data, the SZE survey can determine both $w$ and $\\Omega_m$ to an accuracy of $\\approx 10\\%$ at $3\\sigma$ significance.  Further improvements will be possible from the X--ray survey. The large number of clusters further alleviates the degeneracy between $w$ and both $\\Omega_m$ and $h$, and, as a result, the X--ray sample can determine $w$ to $\\approx 10\\%$ and $\\Omega_m$ to $\\approx 5\\%$ accuracy, in combination with either the CMB or the SN data.  Our work focuses primarily on the statistics of cluster surveys.  We have provided an estimate of the scale of various systematic uncertainties.  Further work is needed to clarify the role of these uncertainties, arising especially from the analytic estimates of the scaling of the mass limits with cosmology, the dependence of the cluster mass function on cosmology, and our neglect of issues such as galaxy formation in the lowest mass clusters.  However, our findings suggest that, in a flat universe, the cluster data lead to tight constraints on a combination of $\\Omega_m$ and $w$, especially valuable because of their high complementarity to those obtained from the CMB anisotropy or Hubble diagrams using SNe as standard candles."
    },
    "0002/astro-ph0002100_arXiv.txt": {
        "abstract": "{ The first sources of ionizing radiation to condense out of the dark and neutral IGM sent ionization fronts sweeping outward through their surroundings, overtaking other condensed objects and photoevaporating them. This feedback effect of universal reionization on cosmic structure formation is demonstrated here for the case of a cosmological minihalo of dark matter and baryons exposed to an external source of ionizing radiation with a quasar-like spectrum, just after the passage of the global ionization front created by the source. We model the pre-ionization minihalo as a truncated, nonsingular isothermal sphere in hydrostatic equilibrium following its collapse out of the expanding background universe and virialization. Results are presented of the first gas dynamical simulations of this process, including radiative transfer. A sample of observational diagnostics is also presented, including the spatially-varying ionization levels of C, N, and O in the flow if a trace of heavy elements is present and the integrated column densities of H~I, He~I and II, and C~IV thru the photoevaporating gas at different velocities which would be measured in absorption against a background source like that responsible for the ionization.  }   \\nonstopmode  \\begin{document} ",
        "introduction": "\\label{sec:intro} Observations of quasar absorption spectra indicate that the universe was reionized prior to redshift $z=5$. Recent detections of H Lyman alpha emission lines from sources at even higher redshift ($z\\leq5.6$) strengthen this conclusion \\cite{Hu98,Wey98}. As new discoveries push the observable horizon back in time ever-closer to the epoch of reionization at redshift $z>5$, increased awareness of its importance as a missing link in the theory of galaxy formation has caused a great renewal of interest in universal reionization. If reionization took place early enough, then Thompson scattering of cosmic microwave background (CMB) photons by free electrons in an ionized intergalactic medium (IGM) would also have left a detectable imprint on the CMB observed today. Recent data on the first Doppler peak in the angular spectrum of CMB anisotropy set limits which imply a reionization redshift $z\\lesssim40$ (model-dependent) \\cite{Gri98}. (For a review and references on reionization work prior to 1996, the reader is referred to Shapiro 1995, while more recent developments are summarized in Haiman \\& Knox 1999.)  A review of the theory of reionization is outside the scope of this brief report. In keeping with the focus of this meeting, I will confine myself to a description of some recent progress on the calculational side of the problem. To solve the full reionization problem we must also solve the problem of how density fluctuations led to galaxy formation and this, in turn, led to secondary energy release in the form of ionizing and dissociating radiation by the stars and quasars formed inside early galaxies, as well as other forms of energy release like supernova explosions, jets and winds, which in turn influenced the future course of galaxy formation. The modern context for this description is the Cold Dark Matter (CDM) model, in which a cold, pressure-free, collisionless gas of dark matter dominates the matter density and structure arises from the gravitational amplification of a scale-free power-spectrum of initially small-amplitude, Gaussian-random-noise primordial density fluctuations, in a hierarchical fashion, with small mass objects collapsing out first and merging together to form larger-mass objects which form later. The calculation of these effects poses an enormous multi-scale computational challenge, involving numerical gas dynamics coupled to gravitational N-body dynamics, which raises the bar of cosmological simulation considerably by adding the requirement that radiative transfer effects be included, as well. To simplify matters, I will henceforth focus on just one of the central challenges of reionization theory, the effect of cosmological ionization fronts. ",
        "conclusions": ""
    },
    "0002/astro-ph0002385_arXiv.txt": {
        "abstract": "Typically $\\sim 5$ percent of the total baryonic mass in luminous elliptical galaxies is in the form of cooled interstellar gas. Although the mass contributed by cooled gas is small relative to the mass of the old stellar system in these galaxies, it is almost certainly concentrated within the optical effective radius where it can influence the local dynamical mass.  However, the mass of cooled gas cannot be confined to very small galactic radii ($r \\lta 0.01r_e$) since its mass would greatly exceed that of known central mass concentrations in giant ellipticals, normally attributed to massive black holes.  We explore the proposition that a population of very low mass, optically dark stars is created from the cooled gas.  For a wide variety of assumed radial distributions for the interstellar cooling, we find that the mass of cooled gas contributes significantly ($\\sim$30 percent) to stellar dynamical mass to light ratios which, as a result, are expected to vary with galactic radius. However, if the stars formed from cooled interstellar gas are optically luminous, their pertubation on the the mass to light ratio of the old stellar population may be reduced. Cooling mass dropout also perturbs the local apparent X-ray surface brightness distribution, often in a positive sense for centrally concentrated cooling.  In general the computed X-ray surface brightness exceeds observed values within $r_e$, suggesting the presence of additional support by magnetic stresses or non-thermal pressure.  The mass of cooled gas inside $r_e$ is sensitive to rate that old stars lose mass ${\\dot M_*}$, but this rate is nearly independent of the initial mass function of the old stellar population. ",
        "introduction": "Perhaps the most perplexing and long-standing problem associated with galactic and cluster cooling flows is the uncertain physical nature and spatial distribution of the gas that cools. The apparent absence of large masses of cooled gas in elliptical galaxies has led some to argue that little or no cooling actually occurs and to postulate some source of heating that offsets the radiative losses in X-ray emission. But the energy required to balance radiative losses is prohibitively large and appropriate heating sources may not be universally available. If the expected radiative cooling actually occurs, two questions arise: (1) What is the nature of the objects that condense from the hot gas? and (2) Where is most of the cooled mass located in the galaxy? Regarding the first question, a variety of physical arguments support the hypothesis, or even the inevitability, of low mass star formation (Fabian, Nulsen \\& Canizares 1982; Thomas 1986; Cowie \\& Binney 1988; Vedder, Trester, \\& Canizares 1988; Sarazin \\& Ashe 1989; Ferland, Fabian \\& Johnstone 1994; Mathews \\& Brighenti 1999). Here we shall address the second question in the context of cooling flows in elliptical galaxies where the known stellar mass and light profiles strongly constrain the spatial distribution of cooled gas.  We adopt the generally accepted hypothesis that only stars of very low mass (e.g. $\\lta 0.1$ $M_{\\odot}$) form in cooling flows (e.g. Ferland, Fabian \\& Johnstone 1994), so that the mass to light ratio of the young stellar population formed from the cooling gas is essentially infinite. In view of the difficulties we encounter with this hypothesis, described below, it seems more likely that the stellar population formed from cooled gas extends to somewhat more massive stars that are optically luminous.  Our gas dynamical models for the evolution of hot interstellar gas in giant ellipticals indicate that the origin of the gas varies with galactic radius. Most of the gas in the inner, optically luminous regions originates from the ejected envelopes of evolving stars; gas in the outer halo is supplied by cosmological secondary infall or tidal acquisitions from neighboring galaxies (Mathews \\& Brighenti 1998b). Circumgalactic gas around massive ellipticals is enriched by Type II supernovae that accompanied early star formation. The variability of circumgalactic gas among luminous ellipticals is responsible for some of the enormous dispersion in X-ray luminosity $L_x$ among ellipticals of similar optical luminosity $L_B$ (Mathews \\& Brighenti 1998a).  Since the hot interstellar gas in a bright elliptical emits observable X-rays, it is clearly losing energy. However, as the gas loses energy it is compressed toward the galactic center by gravitational forces and $Pdv$ work maintains the high temperatures observed, $T \\sim 10^7$ K, producing a galactic cooling flow. The positive interstellar temperature gradients typically observed within a few effective radii are often cited as evidence of radiative cooling in a cooling flow, but this cooling is due instead to the mixing of somewhat cooler, locally virialized gas ejected from stars with hotter gas arriving from larger galactic radii (Mathews \\& Brighenti 1998b; Brighenti \\& Mathews 1998, 1999a). If large entropy fluctuations are present in the hot gas, catastrophic cooling can occur at any radius in the flow. Regions of low entropy (low temperature, high density) radiate more and cool sooner. The amplitude distribution of entropy fluctuations in the interstellar gas determines the radius where cooling mass dropout occurs in the cooling flow. For example, if the entropy in some region in the flow is only slightly less than in the ambient flow, the differential radiative cooling will be small and the region will cool out of the flow at small galactic radii; conversely, localized regions with entropy much lower than the ambient flow cool rapidly and deposit their mass at large radii. Some possible sources of interstellar entropy variations are stellar winds, explosions of Type Ia supernovae, non-uniform SNII heating at early times, and mergers with small, gas-rich galaxies.  The total rate that mass cools and drops out of the flow is closely related to the X-ray luminosity $L_x$. The X-ray luminosity can be approximately expressed as the product of the total cooling rate ${\\dot M}$ and the enthalpy per gram in the hot gas, or $${\\dot M} = \\left({2 \\mu m_p \\over 5 k T}\\right) L_{x,bol} \\approx 2.5 M_{\\odot} {\\rm yr}^{-1}.$$ Here we have used data from the giant Virgo elliptical NGC 4472: $T \\approx 1.3 \\times 10^7$ K; $L_x(0.5 - 4.5 ~{\\rm keV}) = 4.5 \\times 10^{41}$; $L_{x,bol} \\approx 1.6 L_x(0.5 - 4.5 ~{\\rm keV})$. If $L_x$ and $T$ are reasonably constant over the Hubble time, a mass $M_{cg} \\approx 3 \\times 10^{10}$ $M_{\\odot}$ of cold gas is expected to condense from the hot ISM somewhere within NGC 4472. Although this mass is very large, it is only about 4 percent of the total stellar mass in NGC 4472 today. The mass that cools can therefore be ignored if it is widely distributed throughout the galactic volume. However, the central concentration of H$\\alpha$ emission in ellipticals (e.g. Macchetto et al. 1996) suggests that the cooling is concentrated toward the galactic center where the interstellar density is highest and the bulk of the X-ray energy is emitted.  The motivation of this paper is to explore a variety of options for the mass dropout profile of cooled gas in bright ellipticals appropriately constrained by the known radial distributions of total stellar and non-baryonic mass. The radial mass dropout profile of cooled gas cannot be determined from first principles because the distribution and amplitude of the entropy and magnetic fluctuations in the hot gas are unknown and difficult to evaluate from simple physical arguments. Nevertheless, the total mass of cooled gas inside an effective (half-light) radius $r_e$ must be consistent with the mass to light ratio determined from stellar velocities and with the total mass inferred from X-ray observations within $r_e$. Assuming that the stellar mass to light ratio is uniform with radius, the stellar mass $M_*(r)$ and the X-ray mass $M_x(r)$ appear to be in nearly perfect agreement for two bright Virgo ellipticals in the range $0.1r_e \\lta r \\lta r_e$ (Brighenti \\& Mathews 1997a). Because of constraints on the mass distribution of cooled gas provided by X-ray and stellar dynamical observations, galactic cooling flows provide a critical venue for testing the physics of mass deposition in cooling flows.  The mass of cold ($T \\lta 10^4$ K) gas $M_{cg}$ actually observed in ellipticals is many orders of magnitude less than the total cooled mass estimated above. For example, neither HI nor H$_2$ gas has been observed in NGC 4472, only upper limits, $M_{cg} \\lta 10^7$ $M_{\\odot}$ (Bregman, Roberts \\& Giovanelli 1988; Braine, Henkel \\& Wiklind 1988). If stars form they must either be indistinguishable from the old stellar population or non-luminous. We explore here the possibility that most of the cooled gas forms dense baryonic clouds or stars that are dark at optical and radio frequencies. The very low mass stars advanced by Ferland, Fabian \\& Johnstone (1994) satisfy this invisibility criterion, while the star formation models of Mathews \\& Brighenti (1999) indicate that (luminous) stars of mass $\\sim 1 - 2$ $M_{\\odot}$ can form in galactic cooling flows.  X-ray studies indicate significant masses of cold, absorbing gas in cluster and galactic cooling flows (e.g. White et al. 1991; Allen et al. 1993; Fabian et al. 1994; Allen \\& Fabian 1997; Buote 1999), but these results are inconsistent with the absence of radio frequency emission from the cold gas (Braine \\& Dupraz 1994; Donahue \\& Vogt 1997) and should be regarded as controversial until this inconsistency is resolved. Even taken at face value, the total mass of cooled gas implied by the X-ray absorption in cluster cooling flows is typically only a small fraction of the total mass that should have cooled in a Hubble time (Allen \\& Fabian 1997; Wise \\& Sarazin 1999), implying that most of the cooled gas may have formed stars. If cooled gas forms into small stars, these stars will have apogalatica near their point of origin where they will spend most of their orbital time. We shall assume that the gas mass that cools and drops out of the flow contributes optically dark (stellar) mass at the radius where the cooling occurred.  In the following we describe a series of gas dynamical calculations for the evolution of X-ray emitting interstellar gas over the Hubble time and investigate a variety of assumptions about the radial distribution of optically dark cooled gas. To be specific, we compare our models with the well-observed massive elliptical NGC 4472. We find that the mass of cooled gas contributes significantly to dynamical mass to light determinations within $r_e$ based on stellar velocities. ",
        "conclusions": "  \\noindent (1) If the hot interstellar gas cools only in the very center of NGC 4472 for $\\sim 10$ Gyrs, the total accumulated mass there would be $\\gta 10$ times larger than the mass of the central black hole observed. If such large concentrated masses were generally present in bright ellipticals, interstellar gas in $r \\lta 0.1r_e$ would be compressed and heated to temperatures $> 1$ keV. Such hot thermal cores are not generally observed. We conclude that the cooling dropout in massive ellipticals must occur before the gas reaches the galactic center. The hypotheses of distributed mass dropout and low mass star formation were proposed many years ago (Fabian, Nulsen \\& Canizares 1982; Thomas 1986; Cowie \\& Binney 1988; Vedder, Trester, \\& Canizares 1988; Sarazin \\& Ashe 1989; Ferland, Fabian, \\& Johnstone 1994). However, these historical arguments were generally based on the notion that mass dropout would help reduce computed X-ray surface brightness profiles at small projected radii, as required by the observations, but we have shown here that in some cases enhanced dropout at small galactic radii can cause $\\Sigma_x$ to increase, not decrease. (The total bolometric X-ray luminosity $L_x$ should always be lower in distributed cooling models since the hot gas experiences only a fraction of the galactic potential.) The best arguments for distributed cooling dropout are (i) limits on the central black hole mass and (ii) the absence of rotational flattening in X-ray images.  \\noindent (2) We have considered a wide variety of possible mass profiles for the radial deposition of cooled interstellar gas in NGC 4472. The dropout mass is assumed to be optically dark, consistent with the formation of very low mass stars. In every case the (stellar plus dropout) mass to light ratio at $\\sim r_e/3$ significantly exceeds the mass to light ratio determined from stellar velocities. If the dropout mass is optically dark, dynamical mass to light ratios in luminous ellipticals should be substantially enhanced by dark baryonic matter. In this case the true $M/L_B$ for luminous stars may be $\\sim 30$ percent smaller than published values.  \\noindent (3) The excellent agreement between the X-ray and ``stellar'' mass in NGC 4472 shown in Figure 1 (and also NGC 4649) in the range $0.1r_e \\lta r \\lta 1 r_e$ may be a coincidence if our estimates of the cooling flow dropout mass are correct and if this mass is non-luminous.  \\noindent (4) Dynamical mass to light determinations within $r_e$ refer to a superposition of two stellar populations: an old luminous population with a de Vaucouleurs galactic profile and a younger population having a bottom-heavy IMF and a different galactic mass profile. If the younger population is optically dark, the mass to light ratio is not likely to be constant with galactic radius within $r_e$.  \\noindent (5) It is not possible to reduce the total amount of mass deposited from the cooling flow simply by lowering the specific rate of stellar mass loss $\\alpha_*(t) \\equiv (dM_*/dt)/M_*$. For all reasonable power law initial mass functions we show that $\\alpha_* \\propto (M_*/L_B)^{-1}$. For given $L_B$, larger stellar masses $M_*(r)$ must accompany lower values of $\\alpha_*$ so the total amount of mass ejected from stars $\\propto \\int \\alpha_*(t)M_{*} dt$ is nearly independent of the IMF.  \\noindent (6) Among the models we consider, those with centrally concentrated mass dropout perform best in minimizing the overall disagreement with the central $M/L_B$ and the X-ray determined mass in $0.1r_e \\lta r \\lta 1 r_e$. Constant $q$ models in which the dropout is proportional to the local gas emissivity at every radius deposit less mass in $0.1r_e \\lta r \\lta 1 r_e$, but may have gas temperatures that are too low ($q \\gta 4$) or central masses that are too large ($q \\lta 1$).  \\noindent (7) Even in the presence of mass dropout, the computed central interstellar gas densities and X-ray surface brightnesses $\\Sigma_x(R)$ are generally too large. Such deviations would be reduced if the hot gas is partially supported in $r \\lta 0.1r_e$ by magnetic or other non-thermal pressure associated with the extended radio source.  \\noindent (8) If the stellar population formed from cooled interstellar gas extends to intermediate masses, $\\sim 1 - 2$ $M_{\\odot}$, its mass to light ratio may blend with that of the older population. In this case the dropout mass would already be represented in the de Vaucouleurs profile representing the stellar mass distribution in our models. Dynamical determinations of $M/L_B$ would be a weighted mean of the two populations. If the dropout stellar population is luminous, some of the difficulties we have discussed here would be alleviated, but not those regarding $\\Sigma_x(R)$.  \\noindent (9) If the influence of rotation and non-thermal pressure can be understood, high resolution images of the central regions of elliptical galaxies using the {\\it Chandra} (AXAF) satellite may detect the presence of dark baryonic dropout material and an accurate determination of the mass to light ratio of the old stellar population."
    },
    "0002/astro-ph0002499_arXiv.txt": {
        "abstract": "We study the long term evolution of a solar type star that is being disrupted by a super massive ($10^6$~M$_\\odot$) black hole. The evolution is followed from the disruption event, which turns the star into a long thin stream of gas, to the point where some of this gas returns to pericenter and begins its second orbit around the black hole. Following the evolution for this long allows us to determine the amount of mass that is accreted by the black hole. We find that approximately 75\\% of the returning mass is not accreted but instead becomes unbound, following the large compression characterizing the return to pericenter.  The impact of a tidal disruption on the surrounding gas may therefore be like that of two consecutive supernove-type events. ",
        "introduction": "\\label{sec:Intro}  It has long been suggested that supermassive black holes (with a black hole mass $M_\\bh\\sim10^6$~M$_{\\odot}$) in relatively low-luminosity active galactic nuclei (AGNs) can be fed by tidal disruptions of stars that are on nearly radial orbits \\citep[e.g.][]{1978MNRAS.184...87F,1982ApJ...262..120L,1985MNRAS.212...23C,1988Natur.333..523R,1990Sci...247..817R}. The frequency of such events is expected to be of the order of $10^{-4}$~yr$^{-1}$ in a galaxy like M31.  In AGNs with more massive central black holes ($M_\\bh\\ga10^8$~M$_{\\odot}$), stars are ``swallowed'' whole, since the ratio of the event horizon radius to the tidal disruption radius increases with the black hole mass $\\left(\\sim M_\\bh^{2/3}\\right)$.  Although direct observational evidence for the disruption process is lacking, a few observations have been tentatively identified as being associated with disruption events. For example, an ultraviolet flare at the center of the elliptical galaxy NGC~4552, although rather weak, has been suggested to result from such an event \\citep{Renzini95}. Other sudden eruptions that have been interpreted as potentially resulting from tidal disruptions include outbursts in IC 3599, NGC 5905, RXJ 1242.6-1119, RXJ 1624.9+7554, and RXJ 1331.9-3243 \\citep[e.g.][and references therein]{Komossa99b,Komossa99a,Grupe99}. It has also been suggested that the sudden appearance and variability of the double-peaked Balmer lines in the active nucleus NGC~1097 is at least broadly consistent with being produced in a ring resulting from tidal disruption \\citep{1995ApJ...443..617S}. Similarly, it has been proposed that an outburst observed in the Seyfert galaxy NGC~5548 was due to a single star falling into a $\\sim10^7$~M$_{\\odot}$ black hole (\\citealt{Peterson86}; although other interpretations exist, e.g.\\ \\citealt{Terlevich88,Kallman88}). More recent work on the broad emission line region of this galaxy, although interesting in its own right, did not shed any new light on the cause of the variability \\citep{Goad98}.  Several aspects of the problem of stellar encounters with supermassive black holes and of tidal disruption have been examined both analytically and numerically \\citep[e.g.][]{1982ApJ...263..377N,1983ApJ...273..749B,1983A&A...121...97C,Hills88,1992ApJ...385..604K,1994ApJ...422..508K,1992MNRAS.255...92S,1993MNRAS.260..463S,1993ApJ...410L..83L,Marck96}. In addition, the interaction of a white dwarf with a massive black hole has been studied both generally \\citep[e.g.][]{1994ApJ...432..680F} and in the context of gamma-ray bursts \\cite[e.g.][]{1999ApJ...520..650F}.  In the present work, we follow the evolution of a star that is tidally disrupted. In particular, we calculate the properties of the disruption debris to longer times than in some of the previous studies, using a Post-Newtonian, Smooth-Particle-Hydrodynamics (SPH) code.  The numerical method is described in $\\S2$, the results are presented in $\\S3$, and a discussion and conclusions follow. ",
        "conclusions": "\\label{sec:Conc} We have performed a 1PN simulation of a solar mass star being disrupted by a super massive ($10^6$~M$_\\odot$) black hole. The disruption process itself causes about 1/2 of the star's matter to become unbound. We follow this matter up to and beyond the time when it returns to pericenter and starts accreting onto the BH. This in turn enables us to determine the amount of returning mass that is actually accreted. Contrary to previous, more heuristic estimates, we find that only about 25\\% of the returning mass actually gets accreted. The rest becomes unbound, following being heated by the strong compression accompanying the approach to pericenter.  The main consequence of this process is that the maximum accretion rate is a factor of 3 lower than expected from a simple examination of the rate of mass return to pericenter. The accretion rate into the volume resolved by our simulation is also quite constant, at about 1~${\\rm M_\\odot~yr^{-1}}$, for the last 20~days of the simulation as opposed to a power law decay which is expected from the rate of return. When integrated over time, these results show that only about 10\\% of the original star's mass actually gets accreted as opposed to the 65\\% expected if all the bound mass would be finally accreted. This means that the mass involved in the accretion flow, be it in the form of an accretion disk \\citep[e.g.][]{1990ApJ...351...38C}, or spherical accretion \\citep[e.g.][]{Loeb97} is considerably smaller than previously estimated.  Consequently, the duration of the expected `flare' can be shorter by a factor 4--5 compared to these early estimates, which makes the detectability of these events much harder (the rather high temperature of some of the infalling debris also makes the bolometric correction high). Indeed, supernova searches in distant galaxies have failed so far to identify any such event unambiguously (Filippenko 2000, private communication).  The mass of the debris that becomes unbound because of the pericenter compression is comparable to the mass of the unbound debris resulting from the original disruption event. The main difference is that this new debris component has a different orbital distribution. Most notably, the velocity has a much larger component in the direction perpendicular to the orbital plane. This additional component of unbound high velocity debris could produce interesting consequences (e.g.\\ components like Sgr~A East in the Galactic center) when colliding with the interstellar medium surrounding the BH \\citep[e.g.][]{1996ApJ...457L..61K}. In particular, the fact that the mass ejection is more spherically symmetric, makes the event more similar to a normal supernova. Thus, tidal disruption events may produce a quite unique signature (in terms of their impact on surrounding gas), in which two supernova-type events are separated by a few weeks to a few months, with the first one being very anisotropic (mass being ejected within a solid angle $\\Omega\\apgt16(R_*/R_{\\rho})^{1/2}(M_*/M_{\\rm bh})^{1/2}$ rad$^2$) and the second more spherically symmetric."
    },
    "0002/astro-ph0002450_arXiv.txt": {
        "abstract": "A non-stationary polar gap model first proposed by Ruderman \\& Sutherland (1975) is modified and applied to spark--associated pulsar emission at radio wavelengths. It is argued that under physical and geometrical conditions prevailing above pulsar polar cap, highly non-stationary spark discharges do not occur at random positions. Instead, sparks should tend to operate in well determined preferred regions. At any instant the polar cap is populated as densely as possible with a number of two-dimensional sparks with a characteristic dimension as well as a typical distance between adjacent sparks being about the polar gap height. Our model differs, however, markedly from its original ``hollow cone'' version. The key feature is the quasi-central spark driven by $\\gamma-B$ pair production process and anchored to the local pole of a sunspot-like surface magnetic field. This fixed spark prevents the motion of other sparks towards the pole, restricting it to slow circumferential ${\\bf E}\\times{\\bf B}$ drift across the planes of field lines converging at the local pole. We argue that the polar spark constitutes the core pulsar emission, and that the annular rings of drifting sparks contribute to conal components of the pulsar beam. We found that the number of nested cones in the beam of typical pulsar should not excced three; a number also found by Mitra \\& Deshpande (1999) using a completely different analysis. We confront predictions of our model with a variety of pulsar data, including mean profiles morphology and their predicted correlations with properties of the $P-\\dot{P}$ pulsar diagram as well as detailed studies of drifting subpulses (Deshpande \\& Rankin 1999). We demonstrate that, if the observing geometry is known, the average profile as well as the apparent drift pattern are fully determined by the values of $P$ and $\\dot{P}$. In the accompanying Paper II we develop a self-consistent theory of coherent pulsar radio emission based on the modified polar gap model explored in this paper. ",
        "introduction": "Pulsar radiation is believed to originate within a bundle of open magnetic field lines, along which a significant part of a rotation induced potential drop can occur. This potential drop accelerates charged particles supplied by the polar cap area, which is a region of the neutron star surface directly connected to the interstellar medium via magnetic field lines. Two types of polar cap activity models have been proposed so far. In the first type, called the free-flow or stationary models \\citep{saf78,as79,a81}, the charged particles flow freely from the polar cap surface and accelerate within a scale height of about one stellar radius $R\\approx 10^6$~cm, due to the potential drop resulting from the curvature of field lines and/or inertia of outstreaming particles. In the second type, called gap models or non-stationary models (Sturrock 1971; Ruderman \\& Sutherland 1975 - RS75 henceforth; Cheng \\& Ruderman 1977, 1980), the free outflow from the polar cap surface is strongly impeded, which leads to the formation of an empty gap just above the polar cap. The high potential drop across the gap is discharged by the photon-induced pair creation in the strong and curved magnetic field. This breakdown of the polar gap stabilizes its height to about one mean free path for the $\\gamma-B$ pair production process, which has to be smaller than the polar cap radius $r_p\\approx 10^4P^{-1/2}$~cm, where $P$ is pulsar period in seconds. The charges of opposite signs are accelerated in opposite directions to energies exceeding $2m_ec^2$ and a cascade avalanche develops. The polar gap breaks in the form of a number of isolated, short-lived discharge tubes called sparks \\citep{cr77,b82,fr82}, which can be naturally interpreted in terms of subpulse-associated plasma columns, modulating the pulsar radiation spatially on the subpulse time-scales ($\\sim 1$~ms).  The plan of the rest of this is as paper follow. In section 2 we revise the RS75 model in such a way that it is no longer the hollow cone type model. We assume that the surface magnetic field is non-dipolar, with a typical radius of curvature much smaller than the neutron star radius ${\\cal R}<10^6$ \\citep{bah83,rom90,rud91a,k91,a93,cr93,mkb99}. We also assume that the planes of the magnetic fields lines should tend to converge at the local pole, introducing some degree of axial symmetry like in the case of ``sun-spot'' field configurations \\citep[][see also Appendix]{cr93,gm99}. We follow the RS75 idea that the high potential drop above the polar cap is discharged via a number of localized spark filaments and argue that one of them should be anchored to the local surface pole, while others should perform more or less ordered circumferential motion around it due to the well-known ${\\bf E}\\times{\\bf B}$ drift. However, we take into account reduction of the vacuum accelerating potential drop (eq. [\\ref{deltamax}]) due to spark development (eq. [\\ref{calef}]) , which reduces also the apparent drift rate (eqs. [\\ref{vperp}], [\\ref{pehat}] and [\\ref{N}]). We argue that both the characteristic dimension as well as the typical distance between sparks should be about the RS75 gap height (eq. [\\ref{h}]). We use this result in section 3, where we propose that polar spark is associated with the core pulsar emission, while concentric rings formed by drifting sparks correspond to nested conal components of the ``non-hollow cone'' radiation pattern. We calculate the number of cones in the pulsar beam and, since it depends mainly on the values of basic pulsar parameters $P$ and $\\dot{P}$, we examine correlations of number of profile components and profile types with properties of standard $P-\\dot{P}$ diagram. The degree of correlations is high, supporting our picture. In section 4 we revise the RS75 subpulse drift model and apply it to a number of pulsars with drifting subpulses. In section 5 we incorporate the behaviours of drifting subpulses into the core/cone morphological model of pulsar proposed and developed by \\citet{r83}. We discuss both general and specific implications of our model in section 6. In the Appendix we briefly discuss evidence of nondipolar surface magnetic field from PSR J2144-3933, the recently reported 8.5-s pulsar.  Despite a number of well-known problems with the neutron star crust binding energy \\citep[see][for short review]{xqz99}, the recent papers concerning pulsars with drifting subpulses PSR 0943+10 \\citep{dr99} and PSR 0031$-$07 \\citep{vj99} strongly suggest that the RS75-type vacuum gap does exist above the polar cap. Motivated by these evidences, \\citet{xqz99} have recently  argued that the binding energy problem can be completely solved if pulsars (at least those with drifting subpulses) represent bare strange stars, which in many aspects are indistinguishable from neutron  stars.  Sparks proposed in the RS75 model as primary sources of the subpulse-associated plasma columns have been criticized, mainly due to short dynamical time scales ($\\sim 10~\\mu$s) as compared with subpulse time scales ($\\sim 1$~ms). To account for the subpulse emission, they would have to repeat at approximately the same place for a time long relative to their lifetime. No clear explanation for such kind of a surface memory was proposed by RS75. In this paper we propose a modification of the RS75 polar gap model and confront the predictions of the revised model with the much larger body of present day observational data, concerning both subpulses in single pulses and average profiles.  New ingredients of our modified RS75 model include: (a) estimation of a characteristic spark dimension ${\\cal D}\\approx h$ (eqs. [\\ref{h}] and [\\ref{calde}]), (b) introduction of quasi-central spark anchored to the local pole of surface magnetic pole (thus our model is no longer of a hollow-cone type); (c) reduction of gap potential drop corresponding to subpulse emission by a filling factor ${\\cal F}<1$ (eq. [\\ref{calef}]). The critical assumption we need to adopt is an existence of a local magnetic pole (different from the global dipole) at the polar cap. This is necessary for two reasons: (i) a radius of curvature of surface magnetic field is small enough to drive and exponentiate spark discharges in the first place; (ii) the spark fixed at the local pole warrants a persistence of spark arrangement in the form of a quasi-annular pattern leading to nested cone structure of the pulsar beam. Such a multiconal organisation of pulsar beams was first suggested by Rankin (1983, 1993a, b), supported by \\citet{gks93} and confirmed ultimately by \\citet{md99}. The  recently reported pulsar with longest period $P=8.5$~s \\citep{ymj99} suggests that the surface magnetic field in pulsars has a sunspot-like configuration \\citep[][see also Appendix]{cr93,gm99}. We argue that one of the local poles is located at the \\citet{gj69}  polar cap, and assume that pulsar radio emission originates on those multipolar surface field lines which reconnect with dipolar ones in the radio emission region. ",
        "conclusions": "There is a growing evidence from pulsars with orderly drifting subpulses \\citep{dr99,vj99} that the RS75-type polar gap does exist \\citep{xqz99}. Analyzing observational properties of both single pulses and mean profiles we have found that the height scale of inner gap should be close to the value given by RS75. Although the physics behind our model is constrained only by dimensional analysis, it is a useful empirical base for more detailed modeling and interpretation of observations. We demonstrate that all periodicities associated with drifting subpulses are determined mainly by the $P$ and $\\dot{P}$ values. If the observing geometry is known, then one can even reproduce patterns of drifting subpulses. We found amazing agreement between simulated and observed drift-bands in a number of cases (Table 1).  Throughout this paper we assume and/or argue that: \\begin{description} \\item(1) At any given time the polar cap is populated as densely as possible with a number of isolated spark discharges. This number is approximately equal to $a^2\\sim 25\\cdot\\dot{P}_{-15}^{0.6}\\cdot P^{-1.3}$. \\item(2) A spark characteristic dimension as well as the typical distance between sparks is about the polar gap height $h\\approx 3\\cdot 10^3P^{1/7} \\dot{P}_{-15}^{-2/7}$~cm. \\item(3) The life-time (exponentiation time $\\tau\\approx {\\cal R}/c$) of each spark is of the order of 10 microseconds but they tend to reappear in almost the same places on time-scales shorter than $h/c\\ll 10\\mu$sec. \\item(4) The actual surface magnetic field is dominated by multipole components. The planes of field lines tend to converge towards a local pole, which does not coincide in general with the global dipole axis. We believe that the actual surface magnetic field has a ``sunspot-like'' structure with one pole at the dipolar polar cap. \\item(5) One spark is anchored quasi-permanently (oscilating on $10\\mu$s time scale) to the local pole, while other sparks perform a circumferential ${\\bf E}\\times {\\bf B}$ drift around the polar spark. \\item(6) Sparks supply a corresponding subpulse-associated plasma columns that can give rise to coherent radio emission due to some instability developing in a purely dipolar magnetospheric region at altitudes of about few percent of the light-cylinder radius $R_L=cP/2\\pi$. \\item(7) The polar spark is associated with the core pulsar emission while other sparks circulate around the pole and contribute to the conal pulsar emission. \\end{description}  The assumptions and/or arguments used in this paper lead to nested-cones model of the mean pulsar beam. Such a beam structure has already been deduced by Rankin (1983) as a result of her profile classification scheme. This has been confirmed later by the analysis of the period dependence of pulse widths of different profile species (Rankin 1993 a,b; Gil, Kijak \\& Seiradakis 1993; Kramer et al. 1994; Gil \\& Krawczyk 1996, Gil \\& Han 1996, Mitra \\& Deshpande 1999). It is worth mentioning that our arguments are completely different from those of Rankin. Therefore, the fact that our model predictions are in excellent agreement with Rankin's classification scheme is not trivial.  The number of nested cones surrounding the core beam depends in our model mainly on basic pulsar parameters $P$ and $\\dot P$ (a dependence on the unknown parameters characterizing the structure of the actual surface magnetic field is rather weak). Therefore we can make three strong model predictions: \\begin{description} \\item(a) The Core Single ($S_t$) and Conal Single ($S_d$) profile pulsars should be well separated on the $P,\\dot{P}$ diagram (Fig.~6). A division line corresponds to $a\\approx r_p/h\\sim 10$ (yellow dotted). \\item(b) The single pulse emission in $S_t$ pulsars (with $a>10$) should be amorphous with no hints of modulation on the subpulse time scales, while $S_d$ pulsars should demonstrate subpulses in their single pulses. \\item(c) The subpulses in core components of complex profiles should be longitude-stationary while an apparent subpulse drift or periodic intensity modulation should occur exclusively in conal components. \\end{description}  These prediction, which seem to be supported observationaly, are a direct consequence of a non-stationary polar gap discharge through a number of isolated sparks, reappearing in almost the same place for a long time as compared with their lifetime. This is an important advantage of sparking models over stationary free-flow models, where all the above observational features have no natural, self-consistent explanation. Perhaps an even more important advantage of a non-stationary polar gap discharge is the generation of a two-stream plasma instability in the magnetosphere near the neutron star (Usov 1987;  Asseo \\& Melikidze 1998). The sparking phenomenon creates a succession of plasma clouds moving along magnetic field lines, each containing particles with a large spread of momenta. Overlapping of particles with different energies from successive clouds ignites strong Langmuir oscilations, which may lead eventualy to the generation of coherent pulsar radio emission (Melikidze, Gil \\& Pataraya 1999). Interestingly, this instability is the only one which, according to our knowledge, develops at altitudes of the order of 1 percent of the light cylinder radius, where the pulsar radio emission is expected to originate (Cordes 1978, 1992; Kijak \\& Gil 1997, 1998). However, the latter statement is again true only if the potential drop height scale in the actual polar gap is close to that described by the RS75.  In the accompanying paper Melikidze, Gil \\& Pataraya (1999) examine a non-linear evolution of Langmuir electrostatic oscillations in plasma clouds associated with sparks reappearing at approximately the same place of the polar cap, according to a model presented in this paper. They found that a well known modulational instability leads to formation of a ``bunch-like'' charged solitons, capable of generating coherent curvature radiation at radio wavelenghts. Thus, we believe that the non-stationary sparking discharge of the polar gap driven by nondipolar surface magnetic field explains not only observational characteristics of pulsar radiation modulation but also the mysterious generation mechanism of this radiation, all in a self-consistent way.   \\vspace{.5cm} \\ni {\\bf Acknowledgements} This paper is supported in part by the KBN Grant 2~P03D~015~12 of the Polish State Committee for Scientific Research. We thank A.A. Deshpande, J.M. Rankin, M. Vivekanand and B.C. Joshi for providing us raw data with drifting subpulses. We also thank K.S. Cheng, T.H.~Hankins, A. Krawczyk, G. Melikidze, D. Mitra and L. Nowakowski  for their valuable contributions to this work. We thank F.C. Michel, D. Nice, B. Zhang and L. Zhang for helpful discussions and especially Z. Arzoumanian, M. Kramer and D. Lorimer for critical reading of parts of the manuscript and many useful comments. We also thank E. Gil for technical assistance. The Arecibo Observatory is operated by Cornell University under a Cooperative Agreement with the U.S. National Science Foundation.   \\appendix"
    },
    "0002/astro-ph0002406_arXiv.txt": {
        "abstract": "We derive the astrometric orbit of the photo-center of the close pair $\\alpha$ UMi AP (= $\\alpha$ UMi Aa) of the Polaris multiple stellar system. The orbit is based on the spectroscopic orbit of the Cepheid $\\alpha$ UMi A (orbital period of AP: 29.59 years), and on the difference $\\Delta\\mu$ between the quasi-instantaneously measured HIPPARCOS proper motion of Polaris and the long-term-averaged proper motion given by the FK5. There remains an ambiguity in the inclination $i$ of the orbit, since $\\Delta\\mu$ cannot distinguish between a prograde orbit $(i = 50\\,.\\!\\!^\\circ1)$ and a retrograde one $(i = 130\\,.\\!\\!^\\circ2)$. Available photographic observations of Polaris favour strongly the retrograde orbit. For the semi-major axis of the photo-center of AP we find about 29 milliarcsec (mas). For the component P, we estimate a mass of 1.5 ${\\cal M}_\\odot$ and a magnitude difference with respect to the Cepheid of 6.5 mag. The present separation between A and P should be about 160 mas.  We obtain the proper motion of the center-of-mass of $\\alpha$ UMi AP with a mean error of about 0.45 mas/year. Using the derived astrometric orbit, we find the position of the center-of-mass at the epoch 1991.31 with an accuracy of about 3.0 mas. Our ephemerides for the orbital correction, required for going from the position of the center-of-mass to the instantaneous position of the photo-center of AP at an arbitrary epoch, have a typical uncertainty of 5 mas. For epochs which differ from the HIPPARCOS epoch by more than a few years, a prediction for the actual position of Polaris based on our results should be significantly more accurate than using the HIPPARCOS data in a linear prediction, since the HIPPARCOS proper motion contains the instantaneous orbital motion of about 4.9 mas/year = 3.1 km/s. Finally we derive the galactic space motion of Polaris. ",
        "introduction": "\\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics{h2016fg1.eps}} \\caption[] {Astrometric orbit (prograde or retrograde) of the photo-center of $\\alpha$ UMi AP. The retrograde orbit is our preferred solution. For detailed explanations see Sect. 3.2.4\\,.} \\end{figure}      Polaris ($\\alpha$ Ursae Minoris, HR 424, HD 8890, ADS 1477, FK 907, HIP 11767) is a very interesting and important object, both from the astrophysical point of view and from the astrometric one. For astrophysics, the most remarkable feature of the multiple stellar system Polaris is the fact that its main component, namely $\\alpha$ UMi A, is a Cepheid variable with a very unusual behaviour. In astrometry, Polaris is one of the most frequently and accurately observed objects, mainly because it is located so close to the North celestial pole and can be used for calibration purposes.  Up to now, the binary nature of Polaris was essentially neglected in ground-based fundamental astrometry, e.g in the FK5 (Fricke et al. 1988). This was justified by the limited accuracy reached by the meridian-circle observations. Now, the high-precision astrometric measurements carried out with the HIPPARCOS satellite (ESA 1997) require strongly to take into account the binary nature of Polaris in order to obtain an adequate astrometric description of $\\alpha$ UMi. Similar procedures are required for many other binaries among the fundamental stars in order to be included properly into the Sixth Catalogue of Fundamental Stars (FK6; Part I: Wielen et al. 1999c; see also Wielen et al. 1998).  The main purpose of the present paper is to obtain a reliable astrometric orbit for Polaris, and to use this astrometric orbit for obtaining high-precision values for the position and proper motion of Polaris. This is done by combining the known spectroscopic orbit of Polaris with ground-based astrometric data given in the FK5 and with the HIPPARCOS results. Before doing so in the Sects. 3 and 4, we present in Sect. 2 an overview of the Polaris system. ",
        "conclusions": ""
    },
    "0002/astro-ph0002295_arXiv.txt": {
        "abstract": "Recently, there have been numerous analyses of the redshift space power spectrum of rich clusters of galaxies. Some of these analyses indicate a \\lq\\lq bump\\rq\\rq~ in the Abell/ACO cluster power spectrum around $k = 0.05h$Mpc$^{-1}$. Such a feature in the power spectrum excludes most standard formation models and indicates possible periodicity in the distribution of large-scale structure. However, the data used in detecting this peak include clusters with estimated redshifts and/or clusters outside of Abell's (1958) statistical sample, {\\it i.e.} $R = 0$ clusters. Here, we present estimates of the redshift-space power spectrum for a newly expanded sample of 637 $R \\ge 1$ Abell/ACO clusters which has a constant number density to z = 0.10 in the Southern Hemisphere and a nearly constant number density to z = 0.14 in the Northern Hemisphere. The volume sampled, $\\sim 10^8h^{-3}$Mpc$^{3}$, is large enough to accurately calculate the power per mode to scales approaching $10^3h^{-1}$Mpc. We find the shape of the power spectrum is a power-law on scales $0.02 \\le k \\le 0.10h$Mpc$^{-1}$, with enhanced power over less rare clusters such as APM clusters. The power-law here follows $n = -1.4$. The power spectrum is essentially featureless, although we do see a dip near $k = 0.04h$Mpc$^{-1}$ which cannot be considered statistically significant based on this data alone. We do not detect a narrow peak at $k \\sim 0.05h$Mpc$^{-1}$ and there is no evidence for a turn-over in the power spectrum as has been previously reported. We compare the shape of the Abell/ACO rich cluster power spectrum to various linear models. ",
        "introduction": "There has recently been a renewed interest in accurately determining the power spectrum of matter distribution scales greater than 100$h^{-1}$Mpc; in part due to the increased number of clusters with measured redshifts and the large volumes they trace. The power spectrum for the galaxy distribution has been determined many times for many different classes of galaxies. However, most galaxy surveys lack the volume necessary for the accurate quantification of power on large-scales ({\\it e.g.} the Las Campanas Redshift Survey (Lin {\\it et al.} 1996-hereafter LCRS) or the Stromlo-APM survey (Tadros \\& Efstathiou 1996). A summary of the results from these analyses is that the redshift-space power  spectra roughly agree on scales $\\lambda$ ($ = 2\\pi/k$) $< 100 h^{-1}$Mpc. In this region, $P(k) \\propto k^n$ and $n \\sim -2$. (Of course, the amplitude of the power spectra depends on the samples of galaxy examined which provides strong evidence for a luminosity bias (see e.g. Vogeley et al. 1992 and Park et al. 1994)). However, on scales $\\lambda > 100h^{-1}$Mpc, there is much less agreement. For example, some galaxy samples, such as from the  LCRS and the Automated Plate Machine (APM) 2d and 3d surveys show a  broad flattening around $k = 0.05 h$Mpc$^{-1}$ although no distinct maximum can be found within convincing statistical bounds (LCRS; Tadros and Efstathiou 1996; Peacock 1997; Gatzanaga \\& Baugh 1998). However, Landy {\\it et al.} (1996) find a distinct peak in $P(k)$ for a 2 dimensional analysis of the LCRS and Broadhurst {\\it et al.} (1990) find a peak near $\\lambda = 130h^{-1}$Mpc in a deep pencil beam survey.  Some Abell/ACO cluster analyses also show a peak around $ k \\sim 0.05 h$Mpc$^{-1}$ (Retzlaff {\\it et al.} 1998; Einasto {\\it et al.} 1997-hereafter E97). Yet other cluster analyses only show a smooth turnover in the power spectrum to its scale-invariant ($n = 1$) form. For instance, the APM clusters, examined by Tadros, Efstathiou and Dalton (1998-hereafter TED98) show a maximum in $P(k)$ at the smaller value of $k \\sim 0.03h$Mpc$^{-1}$ and no distinct \\lq\\lq bump\\rq\\rq~ at $k = 0.05h$Mpc$^{-1}$. Also, Peacock \\& West (1992) and Jing \\& Valdarnini (1993) find a break in the Abell cluster power spectrum near $k = 0.05h$Mpc$^{-1}$, but no distinct peak in power. An excellent review of the power spectra for different galaxy species can be found in Einasto {\\it et al.} (1999-hereafter E99). E99 determine a mean power spectrum for all galaxies for a large range in wavenumber. They do this by using the APM 2d power spectra on small scales, and by averaging over numerous samples on large scales and then normalizing to the APM 2d power.  If a narrow peak in power near $k \\sim 0.05h$Mpc$^{-1}$ is a real feature of the power spectrum in general, most current models of structure formation (in the quasi-linear regime) become invalid (E99). While baryonic signatures can produce features in the power spectrum, those features are oscillatory and they do not produce a singular, narrow peak as seen in some of the current data. Eisenstein et al. (1998) examined this prospect and found that no selection of cosmological parameters reproduces the power spectrum as seen in E97. However, Gramann \\& Suhhonenko (1999) suggest that an inflationary scenario with a scalar field having a localized step-like feature can reproduce the power spectrum of clusters. However, in this work, we show that the peak in the cluster power spectrum is not present in larger (in volume and in number) cluster samples after excluding less reliable data (such as $R=0$ Abell/ACO clusters and clusters with estimated redshifts).  Our aim in this paper is to provide an estimate for the power spectrum of Abell/ACO clusters that is based on a complete and fair sample. Both Retzlaff {\\it et al.} and E97 use $R =0$ clusters in their determination of $P(k)$. Einasto {\\it et al.} (1994) have argued that $R =0$ clusters do not contaminate studies of large-scale structure because the multiplicity of superclusters is independent of richness and the mean separation distances for $R =0$ and $R \\ge 1$ clusters are very similar. However,  $R = 0$ clusters were not cataloged in a systematic way and were never meant to be examined in a statistical manner due to their incompleteness (Abell 1958). In addition, many researchers have found line-of-sight anisotropies in $R = 0$ cluster samples (Sutherland 1988; Efstathiou {\\it et al.} 1992; Peacock \\& West 1992). Therefore, the use of $R =0$ clusters in the determination of $P(k)$ is highly suspect. E97 have also used a large number (435 out of 1305 clusters) of estimated redshifts in their determination of $P(k)$. We also suspect that E97 used a large number of cluster redshifts with only one measured galaxy. Miller {\\it et al.} (1999a) show that cluster velocities with one measured galaxy are in error by more than $2500$ km s$^{-1}$ 14\\% of the time. Of course, estimated redshifts are only accurate to at best 25\\%. Thus, the statistical certainty of any large-scale structure analyses based on the cluster samples with a large number of estimated or poorly determined redshifts must also be taken with caution. ",
        "conclusions": "The agreement between the shapes of $P(k)$ for the four different samples shown in Figure 6 (from $k = 0.05$ to $0.15 $h Mpc$^{-1}$), provides further evidence that clusters are tracers of the peaks of the underlying luminous mass distribution. While there is a great deal of volume-overlap in these four samples, they are made up of significantly different luminous objects (from very poor APM clusters to the richest Abell clusters). For instance, the Retzlaff et al. (1998) Abell/ACO sample contains at most 253 $R \\ge 1$ clusters, while the remaining 218 are $R \\ge 0$. Our sample contains 637 $R \\ge 1$ clusters. The APM sample of 364 clusters, contains even fewer $R \\ge 1$ Abell clusters ($\\sim 40$). If all groups and clusters trace the underlying mass distribution in a similar way, then the we would expect their respective power spectra to be similar in shape, and only the amplitude to vary.  Previous analyses of the cluster power spectrum have been plagued by three major problems: (1) uncertainties in the number density, (2) small volumes, and (3) irregularly shaped volumes. The sample analyzed in this work greatly improves upon each of these difficulties. Our Abell/ACO sample has a nearly constant number density throughout the entire volume. This is in stark contrast to most other sparse tracer surveys (such as the QDOT {\\it IRAS} survey power spectrum of FKP and the Retzlaff {\\it et al.} Abell/ACO cluster sample). Along with the number density, the large size of the volume and the semi-regular shape of the double-cone geometry, all contribute significantly to a more accurate determination of $P(k)$ on the largest scales. The reality of the power on scales $200 - 300h^{-1}$Mpc is also becoming evident observationally. Batuski {\\it et al.} 1999 have recently discovered two filamentary superclusters in the constellation of Aquarius that are as long as $75h^{-1}$Mpc and $150h^{-1}$Mpc.  As we peer out further into the local Universe, we continue to find structures on very large scales.   We have presented the redshift-space power spectrum for the largest galaxy cluster sample compiled to date. This sample has been examined extensively for projection effects, anisotropies, and observational selection effects and found to be a fair and complete sampling of biased matter in the local Universe. The volume and shape of the survey provide accurate and robust measurements of $P(k)$ over the wavenumber range $0.015 \\le k \\le 0.15 h$Mpc$^{-1}$. From $k= 0.15$ down to $k = 0.05 h$Mpc$^{-1}$, we find a similar shape to the power spectrum compared to other cluster samples such as the APM cluster survey and a smaller sample of $R \\ge 0$ Abell/ACO clusters studied by Retzlaff {\\it et al.} (1998). At smaller $k$, we do not find any statistically significant features in $P(k)$. Unlike previous cluster $P(k)$ analyses, we do not find any strong evidence for a turnover. We find that $\\Lambda$CDM and low $\\Omega_0$ CDM linear models provide excellent fits to the rich cluster power spectrum.  {\\bf Acknowledgments} The authors wish to thank Adrian Melott and Daniel Eisenstein for helpful conversations. We also would like to acknowledge the role of the the referee, Michael S. Vogeley, for his suggestions on improving the original manuscript. We also thank H. Tadros for supplying the APM cluster PS in electronic form. CM was funded in part by NASA-EPSCoR through the Maine Science and Technology Foundation.  \\begin{deluxetable}{cccccc} \\tablenum{1} \\tablewidth{0pt} \\tablecaption{\\bf Sky Coverage} \\tablehead{ \\colhead{$\\ell$ range} & \\colhead{ $b$ range} & \\colhead{ $z$ range} & \\colhead{Number (all z)} & \\colhead{ Number with} & \\colhead{Fraction{\\tablenotemark{a}}} \\nl \\colhead{ } & \\colhead{ } & \\colhead{ } & \\colhead{cataloged} & \\colhead{observed redshifts} & \\colhead{}} \\startdata $0^o \\le \\ell < 180^o$ & $30^o \\le b \\le 90^o$ & $z \\le 0.10$ & 636 & 80 & 0.1257 \\nl $0^o \\le \\ell < 180^o$ & $30^o \\le b \\le 90^o$ & $0.10 < z \\le 0.14$ & 636 & 86 & 0.1352 \\nl $180^o \\le \\ell < 360^o$ & $30^o \\le b \\le 90^o$ &$z \\le 0.10$ & 503 & 78 & 0.1550 \\nl $180^o \\le \\ell < 360^o$ & $30^o \\le b \\le 90^o$ &$0.10 < z \\le 0.14$ & 503 & 52 & 0.1034 \\nl $0^o \\le \\ell < 180^o$ & $-90^o \\le b \\le -30^o$ &$z \\le 0.10$ &608 & 95 & 0.1563 \\nl $0^o \\le \\ell < 180^o$ & $-90^o \\le b \\le -30^o$ &$0.10 < z \\le 0.14$ & 608 & 84 & 0.1382 \\nl $180^o \\le \\ell < 360^o$ & $-90^o \\le b \\le -30^o$ &$z \\le 0.10$ & 492 & 75 & 0.1524 \\nl \\tablenotetext{a}{Fraction is the Number observed/ Number cataloged.} \\enddata \\end{deluxetable}   \\begin{deluxetable}{ccc} \\tablenum{2} \\tablewidth{0pt} \\tablecaption{\\bf Goodness-of-Fit to Linear Models} \\tablehead{ \\colhead{Model} & \\colhead{$\\chi_{reduced}^2$} & \\colhead{DOF}} \\startdata $\\Lambda$CDM ($H_o = 100$km s$^{-1}$) &  0.65 & 8 \\nl Open CDM ($H_o = 100$km s$^{-1}$) & 0.66 & 8 \\nl CDM ($H_o = 100$km s$^{-1}$) & 4.55 & 8 \\nl MDM ($H_o = 50$km s$^{-1}$) & 2.19 &  8 \\nl \\enddata \\end{deluxetable}"
    },
    "0002/astro-ph0002010_arXiv.txt": {
        "abstract": "We calculate the radiation spectra from advection-dominated accretion flows (ADAFs), taking into account the effects of a global magnetic field. Calculation is based on the analytic model for magnetized ADAFs proposed by Kaburaki, where a large-scale magnetic field controls the accretion process. Adjusting a few parameters, we find that our model can well reproduce the observed spectrum of Sagittarius A$^{*}$. The result is discussed in comparison with those of  well-known ADAF models, where the turbulent viscosity controls the accretion process. ",
        "introduction": "\\footnotetext[1]{Present adress: Department of Earth and Space Science, Graduate School of Science, Osaka University, Toyonaka, Osaka 560-0043, Japan\\\\ kino@vega.ess.sci.osaka-u.ac.jp} The optically thin, advection-dominated accretion flows (ADAFs) have been studied by a number of authors during past several years (e.g. Narayan $\\&$ Yi 1994, 1995a, b; Abramowicz et al.\\ 1995; Nakamura et al.\\ 1997, Manmoto, Mineshige \\& Kusunose 1997; Narayan et al.\\ 1998). These models are very successful in describing both spectra and dynamics of accreting black hole systems such as those in binaries and in low-luminosity active galactic nuclei (AGNs).  The observed spectra can be explained as follows.  The radio emission is due to the synchrotron emission in turbulent magnetic fields in the accretion flow.  These synchrotron photons serve as seed photons for the inverse Compton process by hot electrons. Once-scattered Compton photons are mainly distributed in the optical band and twice-scattered Compton photons, in soft X-ray band. Bremsstrahlung due to electron-electron and electron-proton collisions gives rise to the observed hard X-ray spectra.  Thus, these ADAF models provide a good framework for understanding the observed spectra. In these models,  both angular momentum transfer and energy dissipation in the accretion flow is assumed to be undertaken by the turbulent viscosity whose size is specified by so-called $\\alpha$ parameter. For this reason, hereafter we call this type of models the ``viscous'' ADAF model in this paper. The magnetic fields are regarded as of turbulence origin and are described by another parameter $\\beta$ which specifies the ratio of the magnetic pressure to the gas pressure.  However there is no reason to believe that the turbulent viscosity is the only candidate that controls the accretion processes. Rather, it is quite natural to think that some types of global magnetic fields may play an essential role. Indeed, there are some evidences for the presence of such an ordered magnetic field in the central region of our Galaxy (e.g., Yusef-Zadeh, Morris \\& Chance 1984). As Kato, Fukue \\& Mineshige (1998) has pointed out the hydromagnetic turbulence in accretion disks may also generate global magnetic fields by dynamo processes due to the presence of helical motions. In view of such circumstances, another type of ADAF model has been proposed by one of the present authors (Kaburaki 1999, 2000; hereafter referred to as K99 and K00). In order to distinguish it from the above viscous ADAF models, hereafter we call it the ``resistive''  ADAF models since energy dissipation in the accretion flow is due to the electric resistivity and angular momentum transfer is supported not by the viscosity but by the magnetic stress of a large scale magnetic field.  The purpose of the present study is to calculate the expected radiation spectra from ADAFs in a global magnetic field based on the resistive ADAF model, in order to compare its predictions with those of the viscous ADAF models. As a most suitable candidate for such a comparison, Sgr A$^*$ is taken up here because it has been observed in many wave lengths as the nearest galactic nucleus and its spectrum has been reproduced many times by the successively advancing viscous ADAF models.  In \\S \\ref{scaled}, we introduce the set of analytic solutions for resistive ADAFs in a suitably scaled form and discuss their basic characteristics. The relevant radiation mechanisms and the methods of calculation of the fluxes are described in \\S \\ref{cal}. These schemes are applied to Sgr A$^*$ in \\S \\ref{result} and the results are discussed in comparison with those of the viscous ADAF models. Finally in \\S \\ref{sum}, we summerize the main results and discuss some related issues. ",
        "conclusions": "\\label{sum}  To summarize the examinations in the previous section, both viscous and resistive ADAF models can explain the observed spectrum of Sgr A$^*$ equally well. In spite of large differences in the basic mechanisms working in both models, the calculated spectra are quite similar, except for the extended-disk case in the resistive model. This fact suggests that also the resistive ADAF model is quite powerful in explaining the behavior of other low luminosity AGNs (Narayan, Mahadevan \\& Quataert 1998). In addition to these analogous aspects, the resistive model seems to have a possibility to explain such an essentially different situation as appeared in the extended disk case. In any case, when the presence of an ordered magnetic field should be taken seriously in some AGNs or in some stellar-size black holes then the resistive ADAF model, whose predictions on the radiation spectra are examined in this paper, will serve the purpose.  One of the most remarkable features of the ADAF models is that the mass of the central black hole seems to be determined only from the fitting to the self-absorbed part of the observed spectrum. In the case of Sgr A$^*$, the resistive ADAF model (hereafter restricting to the case of compact disk) predicts the central mass of $3.9\\times 10^5\\ M_{\\odot}$ while the viscous models predict $1.0\\times10^6\\ M_{\\odot}$ (Manmoto et.al 1997) and $2.5\\times 10^6\\ M_{\\odot}$ (Narayan et al. 1998). The accuracy of the fittings for other disk parameters than the black hole mass will be greatly improved by the precise determination of the position and height of the synchrotron peak from observations.  The black hole mass predicted by the resistive ADAF model is evidently smaller compared with the predictions of the viscous ADAF models. The latter values are consistent with the dynamically reduced value of $2.5\\times10^6\\ M_{\\odot}$ (Haller et al.\\ 1996; Eckart \\& Genzel 1997), which may be considered as an upper limit for the black hole mass. In the history of viscous ADAF models, the predicted black-hole mass was as small as $7\\times10^5\\ M_{\\odot}$ (Narayan et al.\\ 1995). Afterwards, by the inclusion of compressive heating, it becomes consistent with the dynamical mass. Since this change is mainly due to the decrease in electron temperature (Narayan et al.\\ 1998), the prediction of the resistive ADAF model may also be increased if the development of its two-temperature versions results in a lower electron temperature. As for the compressive heating, it is already included in the resistive model.  In spite of the resemblance in the predicted spectral shape, there are of course many differences in the predictions of the viscous and resistive models. The precise dependences on the relevant quantities of the luminosities of the self-absorbed part, the synchrotron peak and the X-ray bump are different. Especially, the dependence of the synchrotron peak-frequency on the strength of magnetic field is much stronger for the resistive model. The essential difference in the geometry of an accretion flow may be in the radius of the inner edge rather than in its vertical thickness. The prediction of the resistive ADAF model for the inner-edge radius of the disk around Sgr A$^*$ is $\\sim 20R_{\\rm ms}$, instead of the radius of marginally stable circular orbit $R_{\\rm ms}$. Although this result justifies the neglection of the general relativistic effects in our treatment, various questions may be raised about the behavior of infalling plasmas. As for this point we only present an idea below from a viewpoint of global consistency, because its detailed analyses are beyond the scope of this paper.  Fig.\\ 1 shows an overview of the flow and magnetic field configurations (see K00, for more details). The accretion flow would be decelerated near the inner edge by the presence of a strong poloidal magnetic field which is maintained by the sweeping effect of the flow. As a result, a certain fraction of the accreting plasma will be turned its direction to go along the poloidal field lines, although the remaining fraction may fall into the central black hole. If the poloidal current driven in the accretion disk can close its circuit successfully around distant regions and along the polar axis, a set of bipolar jets will be formed (Kaburaki $\\&$ Itoh 1987). Even if the mechanism for formation of jet does not work well, the plasma within the inner edge is likely to extend to the polar regions.  The presence of the plasma within the inner edge of an accretion disk and near the polar axis can be a possible source of the excess above the self-absorbed slope in radio band of the observed spectrum. Very recent VLBI observations of Sgr A$^*$ (Krichbaum et al. 1998; Lo et al. 1998) report that its intrinsic sizes in the east-west direction at 215 GHz and 68GHz are about 20 $R_{\\rm G}$ (with $M=2.5\\times 10^6M_{\\odot}$). A half of this size (i.e., its radius) is just comparable to the size of the inner edge $\\sim 60R_{\\rm G}$ of our model fitting with $M=4\\times10^5M_{\\odot}$. However, the value of the black hole mass estimated from the spectral fitting may be increased if there is a possibility for Sgr A$^*$ to have a wind-type mass loss from the surfaces of its disk (such possibilities have been noted for various types of objects by, e.g., Blandford \\& Begelman 1999; Di Matteo et al. 1999; Quataert \\& Narayan 1999). In such a case, the VLBI component becomes smaller than the size of the inner edge.  From the standpoint of the resistive ADAF, therefore, the above observations should be interpreted as suggesting the presence of a compact structure which is comparable to or smaller than the inner-edge radius of the accretion disk. In view of the vertical elongation of this component reported by Lo et al. (1998), this structure is very likely to be the root of a jet as suggested by them. This picture is very consistent with the view described above in relation to Fig.\\ 1. In this case, however, the location of the self-absorbed slope in the $\\nu$-$\\nu L_{\\nu}$ diagram should be slightly shifted towards the higher-frequency side so that the VLBI data points can be regarded as an excess from the disk's contribution. \\\\  One of the authors (M.\\ K.) would like to thank Tadahiro Manmoto for many valuable comments on the viscous ADAF models. He is also grateful to Umin Lee for his suggestions on some numerical technics."
    },
    "0002/astro-ph0002360_arXiv.txt": {
        "abstract": "Recent cosmological observations strongly suggest that the universe is dominated by an unknown form of energy with negative pressure. Why is this dark energy density of order the critical density today? We propose that the dark energy has periodically dominated in the past so that its preponderance today is natural. We illustrate this paradigm with a model potential and show that its predictions are consistent with all observations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002156_arXiv.txt": {
        "abstract": "The basic workings of inflationary models are summarized, along with the arguments that strongly suggest that our universe is the product of inflation.  The mechanisms that lead to eternal inflation in both new and chaotic models are described.  Although the infinity of pocket universes produced by eternal inflation are unobservable, it is argued that eternal inflation has real consequences in terms of the way that predictions are extracted from theoretical models.  The ambiguities in defining probabilities in eternally inflating spacetimes are reviewed, with emphasis on the youngness paradox that results from a synchronous gauge regularization technique.  Vilenkin's proposal for avoiding these problems is also discussed. ",
        "introduction": "There are many fascinating issues associated with eternal inflation, so I can think of no subject more appropriate to discuss in a volume commemorating David Schramm.  The shock of Dave's untimely death showed that even the most vibrant of human lives is not eternal, but his continued influence on our entire field proves that in many ways David Schramm is truly eternal. Dave is largely responsible for creating the interface between particle physics and cosmology, and is very much responsible for cementing together the community in which this interface developed.  His warmth, his enthusiasm, and the efforts that he made to welcome young scientists to the field have strengthened our community in a way that will not be forgotten.  I will begin by summarizing the basics of inflation, including a discussion of how inflation works, and why many of us believe that our universe almost certainly evolved through some form of inflation.  This material is not new, but I think it should certainly be included in any volume that attempts to summarize the important advances that Dave helped to develop and promote. Then I will move on to discuss eternal inflation, attempting to emphasize that this topic has important implications, and raises important questions, which should not be dismissed as being metaphysical. ",
        "conclusions": "In this paper I have summarized the workings of inflation, and the arguments that strongly suggest that our universe is the product of inflation.  I argued that inflation can explain the size, the Hubble expansion, the homogeneity, the isotropy, and the flatness of our universe, as well as the absence of magnetic monopoles, and even the characteristics of the nonuniformities. The detailed observations of the cosmic background radiation anisotropies continue to fall in line with inflationary expectations, and the evidence for an accelerating universe fits well with the inflationary preference for a flat universe.  Next I turned to the question of eternal inflation, claiming that essentially all inflationary models are eternal. In my opinion this makes inflation very robust: if it starts anywhere, at any time in all of eternity, it produces an infinite number of pocket universes.  Eternal inflation has the very attractive feature, from my point of view, that it offers the possibility of allowing unique predictions even if the underlying string theory does not have a unique vacuum.  I have also emphasized, however, that there are important problems in understanding the implications of eternal inflation.  First, there is the problem that we do not know how to treat the situation in which the scalar field climbs upward to the Planck energy scale.  Second, the definition of probabilities in an eternally inflating spacetime is not yet a closed issue, although important progress has been made.  And third, I might add that the entire present approach is at best semiclassical.  A better treatment may not be possible until we have a much better handle on quantum gravity, but eventually this issue will have to be faced."
    },
    "0002/astro-ph0002468_arXiv.txt": {
        "abstract": " ",
        "introduction": "Near-infrared luminosity provides a good measure of a galaxy's mass, over a wide range of Hubble types \\cite{Gavazzi}, redshifts, and star formation histories \\cite{Charlot},\\cite{KC}.  With the availability of sensitive, large-format infrared array cameras on large telescopes, it is now practical to obtain infrared galaxy samples reaching below $L^*$ at $z > 1$ over areas large enough to encompass hundreds of such galaxies. In late 2001, the Space Infrared Telescope Facility (SIRTF \\cite{Fanson}) will provide $\\mu$Jy-level sensitivity in the mid-infrared, enabling rest-frame $2\\ \\mu$m-selected samples reaching $L^*$ at $z>3$ to be obtained \\cite{Fazio}.  Kauffmann \\& Charlot\\cite{KC} have recently proposed that the fraction of $z > 1$ galaxies in deep infrared $K$-selected samples provides a powerful means of discriminating between pure luminosity evolution (PLE) and hierarchical (CDM) scenarios for massive galaxy formation and evolution.  They argue that the paucity of $z > 1$ galaxies in the Hawaii $K$-band samples \\cite{Songaila}, \\cite{Cowie} already provides strong evidence against the PLE scenario and is consistent with $\\Omega=1$ CDM models.  Similar arguments have been made based on the absence of red objects in surveys covering several square arcminutes to $K \\sim 22$ \\cite{Zepf}, primarily the KPNO 4m Infrared Imager observations of the Hubble Deep Field (HDF IRIM) obtained by Dickinson et al.  The HDF spans a volume small enough that it would be expected to contain only a few dozen L* galaxies with $1 < z < 2$ \\cite{Dickinson}.  In combination with the strong clustering seen in Lyman break galaxies over substantially larger fields \\cite{Adelberger}, this suggests that global conclusions drawn from samples like the HDF should be treated with some caution.   Using the KPNO 4m, Elston, Eisenhardt, and Stanford (hereafter EES) have recently completed a substantially larger $K$-selected survey, whose properties are summarized and compared to the Hawaii and HDF IRIM surveys in Table 1.  The EES survey is divided into 4 regions around the sky, providing some indication of field-to-field variations caused by clustering. Figure 1 shows a color--color diagram for the survey.  \\medskip \\begin{center} {\\bf Table 1.} IR Field Survey Characteristics. \\end{center} \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|} \\hline Survey & Bands & $K(10\\sigma)$ & Area & N$_{gal}$ & N/A \\\\ &       &  mag & sq arcmin &  & \\#/sq arcmin \\\\ \\hline \\hline Hawaii\\cite{Cowie}    & $B,I,K$ & 19.3 & 26 & 122 & 5 \\\\ \\hline EES & $B,R,I,Z,J,K$ & 20 & 124 & 1683 & 14 \\\\ &               & 19 &     & 720  & 6 \\\\ \\hline HDF-IRIM & $J,H,K$ & 21.2 & 7 & 149 & 21 \\\\ &        & 20   &   & 76  & 11 \\\\ &        & 19   &   & 44  & 6 \\\\ \\hline \\end{tabular} \\end{center}  \\begin{figure} \\plotfiddle{eisenhardt-fig1.ps}{3.0truein}{270}{50}{50}{-180}{270}  \\caption{The $R-J$, $J-K$ color-color diagram for the EES survey. Overlaid on the data are the tracks of SEDs associated with zero redshift galaxies of the indicated types.  Iso-redshift contours are plotted at redshift intervals of 0.5 from redshifts of 0 to 3.5.    Also plotted on the figure are the observed $RJK$ colors of dwarf stars from Leggett (1992). }  \\end{figure} ",
        "conclusions": ""
    },
    "0002/astro-ph0002374_arXiv.txt": {
        "abstract": "We present far-UV observations of the Crab nebula and its pulsar made with the Space Telescope Imaging Spectrograph onboard the {\\it Hubble Space Telescope}. Broad, blueshifted absorption arising in the nebula is seen in C~IV~$\\lambda$1550, reaching a blueward velocity of $\\sim 2500 \\kms$. This can be interpreted as evidence for a fast outer shell surrounding the Crab nebula, and we adopt a spherically symmetric model to constrain the properties of such a shell. From the line profile we find that the density appears to decrease outward in the shell. A likely lower limit to the shell mass is $\\sim 0.3\\Msun$ with an accompanying kinetic energy of $\\sim 1.5\\EE{49}$~ergs. A fast massive shell with $10^{51}$~ergs cannot be excluded, but is less likely if the density profile is much steeper than $\\rho(R) \\propto R^{-4}$ and the maximum velocity is $\\lesssim 6000 \\kms$. The observations cover the region $1140-1720$ \\AA, which is further into the ultraviolet than has previously been obtained for the pulsar. With the time-tag mode of the spectrograph we obtain the pulse profile in this spectral regime. The profile is similar to that previously obtained by us in the near-UV, although the primary peak is marginally narrower. Together with the near-UV data, and new optical data from the {\\it Nordic Optical Telescope}, our spectrum of the Crab pulsar covers the entire region from $1140 - 9250$~\\AA. Dereddening the spectrum with a standard extinction curve we achieve a flat spectrum for the reddening parameters $E(B-V)=0.52$, $R=3.1$. This dereddened spectrum of the Crab pulsar can be fitted by a power law with spectral index $\\alpha_{\\nu} = 0.11\\pm0.04$. The main uncertainty in determining the spectral index is the amount and characteristics of the interstellar reddening, and we have investigated the dependence of $\\alpha_{\\nu}$ on $E(B-V)$ and $R$. In the extended emission covered by our $25 \\arcsec \\times 0\\farcs5$ slit in the far-UV, we detect C~IV~$\\lambda$1550 and He~II~$\\lambda$1640 emission lines from the Crab nebula. Several interstellar absorption lines are detected along the line of sight to the pulsar. The Ly$\\alpha$ absorption indicates a column density of $(3.0\\pm0.5)\\EE{21}$ cm$^{-2}$ of neutral hydrogen, which agrees well with our estimate of $E(B-V)$=0.52 mag. Other lines show no evidence of severe depletion of metals in atomic gas. ",
        "introduction": "The Crab nebula and its pulsar (PSR 0531+21) are among the most studied objects in the sky. The discovery of the Crab pulsar as a fast rotating radio pulsar (\\cite{SR68};~\\cite{C69}) paved the way for the interpretation of pulsars as neutron stars (\\cite{G68}). Also, the position of the Crab pulsar in the center of the Crab nebula, which is the remnant of supernova 1054, clearly supports the supernova - neutron star connection. Soon after the radio detection the pulsar was also shown to emit optical pulsations (Cocke, Disney, \\& Taylor 1969). This established that the pulsating star was the well known south preceding star in the center of the nebula, which early optical spectroscopy showed to emit a featureless continuum (\\cite{M42}). To date, more than 1000 radio pulsars are known, but only the following few have optical counterparts known to pulsate also in visible light: the Crab pulsar (\\cite{Co69}), the LMC pulsar 0540-69 (\\cite{MP85}), the Vela pulsar (\\cite{Wa77}), PSR 0656+14 (\\cite{Sh97}) and (possibly) the Geminga pulsar (\\cite{Sh98}). In the near-UV (NUV), pulsations have only been established for the Crab pulsar (\\cite{P93};~\\cite{G98}, henceforth G98). Due to the faintness of these objects in the optical and in the ultraviolet, the spectroscopic information is very limited. PSR 0540-69 was observed with the Faint Object Spectrograph (FOS) onboard {\\it Hubble Space Telescope (HST)} in the $2500-5000$~\\AA\\ range (\\cite{H97}) and showed a rather steep power law spectrum. These observations were, however, contaminated by nebular emission. The Geminga pulsar was observed with the {\\it Keck} telescope (Martin, Halpern, \\& Schiminovich 1998), but the spectrum has very low signal-to-noise because this pulsar is exceedingly faint in the optical. The only pulsar for which good signal-to-noise spectroscopy in the optical and ultraviolet can be obtained is the Crab pulsar. Surprisingly enough, very little has been done in this respect since the optical observations of Oke (1969). In particular, until the study by G98, no UV spectroscopy of the Crab pulsar had been published since the first attempts by the {\\it International Ultraviolet Explorer (IUE)} (\\cite{Ben80}). The {\\it IUE} data cover only the NUV region ($2000-3150$~\\AA) and have poor signal-to-noise. The {\\it HST}/STIS (Space Telescope Imaging Spectrograph) data from G98, and the new data presented here, clearly supersede these early attempts.  The Crab pulsar has been extensively studied over a very broad wavelength range, from the radio up to $\\gamma$-rays (e.g., \\cite{LGS98}). The high energy emission, from infrared (IR) to $\\gamma$-rays, is believed to be the result of the same emission mechanism (e.g., \\cite{LGS98}). It is therefore of interest to fill in the gaps in the observed spectrum of the pulsar in this range. Although the pulsar is relatively bright in the optical, UV observations are difficult due to the large extinction toward the Crab, $E(B-V)\\sim0.5$ mag (e.g., \\cite{DF85}, and references therein). Here, we present UV observations of the Crab pulsar further into the far-UV (FUV) ($1140-1720$ \\AA) than has previously been obtained. These are presented together with our previous NUV-data ($1600-3200$ \\AA)(G98) and new optical data from the {\\it Nordic Optical Telescope (NOT)}. Due to the large extinction correction, great care must be taken to draw conclusions about the intrinsic spectrum, and thus the emission mechanism of the pulsar. However, this procedure might also give a hint on the absorption properties of the dust in the direction toward the pulsar.   In addition to the pulsar emission, we detect emission lines from the Crab nebula itself in the FUV. In particular, the strength of the C~IV~$\\lambda$1550 emission can be of interest for abundance determinations. Even more interesting is the broad C~IV~$\\lambda$1550 absorption line from the nebula detected against the pulsar continuum. This line provides information on the nature of the SN 1054 event.  Although the Crab nebula has been studied extensively, the nature of the progenitor remains unknown. According to models based on the existence of the central neutron star, as well as on nucleosynthesis arguments, the zero-age main sequence (ZAMS) mass of the progenitor was probably in the range $8-13 \\Msun$ (Nomoto 1985). The amount of material observed in the nebula ($4.6\\pm1.8 \\Msun$) seems too low to account for this (Fesen, Shull, \\& Hurford 1997). Furthermore, the velocities of the filaments ($\\sim 1400 \\kms$, Davidson \\& Fesen 1985) give an uncomfortably low kinetic energy ($\\lesssim 1\\EE{50}$ ergs) compared to other supernova remnants, i.e., at least an order of magnitude less than the canonical energy of supernovae, $10^{51}$ ergs.  The ``missing mass'' could either be in a slow progenitor wind, or in a fast, hitherto undetected, shell  ejected at the explosion (Chevalier 1977). If the latter is true, as is hinted by the observations of the outer [O~III] skin of the nebula (Hester et al. 1996; Sankrit \\& Hester 1997), this shell might account for the missing mass and kinetic energy of the nebula. The question remains, however, why such a shell has escaped detection despite many efforts to observe it (see, e.g., \\cite{Fe97}). One possibility is that the low density of the surrounding gas is not high enough to give rise to detectable circumstellar emission when interacting with the ejecta; neither X-ray nor radio searches have indicated any evidence of circumstellar interaction between fast ejecta and ambient gas (Mauche \\& Gorenstein 1989; Frail et al. 1995). If a fast shell is absent, the birth of the Crab was definitely a low energy event. This would call for a revision of our understanding of supernova explosions, especially since SN 1054 apparently was not unusually dim according to historical records (Chevalier 1977; Wheeler 1978).  It is thus of great interest to further investigate whether there is a stellar wind or supernova ejecta outside the observed nebula, and what velocity this gas may have.  Lundqvist, Fransson, \\& Chevalier (1986, henceforth LFC86) proposed to search for a fast shell by looking in the UV toward the Crab pulsar. Their time dependent photoionization calculations showed that C~IV~$\\lambda$1550 could show up in blueshifted absorption if the ionization history of the shell was as predicted in some models of Reynolds \\& Chevalier (1984). Here, we present the detection of this broad absorption line, and discuss its implications for the fast shell around the Crab nebula.  First we discuss the observations and reductions (\\S 2). We then (\\S 3) discuss the pulse profile, the amount of reddening toward the pulsar and the intrinsic pulsar spectrum in the optical/UV. In \\S 3 we also discuss the lines originating from the interstellar gas toward the pulsar and from the Crab nebula itself. Some of these observations constrain the properties of a possible outer shell. In \\S 4 we summarize our conclusions. ",
        "conclusions": "Using STIS onboard the {\\it HST} we have observed the Crab nebula and its pulsar in the far-UV ($1140-1720$ \\AA). We have obtained the pulse profile of the pulsar, which is very similar to our previous near-UV profile, although the primary peak appears to be marginally narrower than in the near-UV data ($5\\%,2\\sigma$). Combining the far- and near-UV data, and assuming an intrinsic power law for the pulsar continuum, we have derived an extinction of $E(B-V)=0.52$ mag toward the Crab. No evidence for a non-standard extinction curve was found. We have also added optical spectra taken with the {\\it NOT} to obtain a spectrum of the pulsar from 1140 \\AA~to 9250 \\AA. We have shown that the pulsar spectrum can be well fitted over the full UV/optical range by a power law with spectral index $\\alpha_{\\nu}=0.11$. The exact value of the spectral index is, however, sensitive to the amount and characteristics of the interstellar reddening, and we have investigated this dependence for a likely range of $E(B-V)$ and $R$. In the optical, we find no evidence for the dip in the pulsar spectrum around 5900 \\AA\\ reported by Nasuti et al. (1996).  The interstellar absorption lines detected in the UV have been analyzed, and are consistent with normal interstellar abundances. The column density of neutral hydrogen is $(3.0\\pm0.5)\\EE{21}$~cm$^{-2}$, which corresponds well to the value derived for $E(B-V)$. From the Crab nebula itself we detect the emission lines C~IV~$\\lambda$1550 and He~II~$\\lambda$1640. The ratio of the fluxes of these lines is similar to what has been derived previously, although obtained with much improved spatial resolution.  C~IV~$\\lambda$1550 is also seen in absorption toward the pulsar. The line is broad and blueshifted with a maximum velocity of  $\\sim 2500 \\kms$, and there is no absorption at zero velocity. These are the highest velocities measured in the Crab and shows that there exists material outside the visible nebula. This can be interpreted as evidence for the fast shell that has been predicted to surround the Crab nebula (Chevalier 1977). We have used a simple, spherically symmetric model in which the density in the shell falls off with radius as $R^{-\\eta}$ from $5\\EE{18}$~cm (corresponding to $\\approx 1680 \\kms$) to derive the mass and energy of such a shell. The conclusions from our model depend on how we tie our model into other observations and models. From the C~IV line alone, we find that the minimum mass and kinetic energy of the fast gas are $\\sim 0.3 \\Msun$ and $\\sim 1.5\\EE{49}$~ergs, respectively. This occurs for a density slope $\\eta \\sim 3$. A model with a flat ($\\eta = 0$) density profile appears unlikely as the required ionization structure disagrees with the modeling of Lundqvist et al. (1986). The maximum mass of the shell is set by progenitor models, and is unlikely to be much larger than $8 \\Msun$. With a high shell mass, and the velocity extending to velocities much higher than we can detect, the shell could carry an energy of $10^{51}$~ergs. The signal-to-noise of C~IV~$\\lambda$1550 is too low at high velocities to reject or confirm such a conclusion.  Adding constraints from the model of Sankrit \\& Hester (1997) to those from the C~IV line narrows down the parameter space for the shell. In particular, a density slope of $\\eta \\gtrsim 3$ is required to agree with the interpretation of the observed [O III]-skin being a radiative shock. For $\\eta \\leq 9$, the shell mass is then $\\gtrsim 1.7 \\Msun$ and the kinetic energy $\\gtrsim 7\\EE{49}$~ergs. Although the limit on the H$\\alpha$ surface brightness from the search of Fesen et al. (1997) tend to favor models with steep density profiles, a model with $\\eta = 3$ might still be possible, if the He/H ratio is higher than solar also in the fast shell, and the asymmetry of the outer shell different from that in the model of Sankrit \\& Hester (1997)."
    },
    "0002/astro-ph0002004_arXiv.txt": {
        "abstract": "When some magnetic field lines connect a Kerr black hole with a disk rotating around it, energy and angular momentum are transferred between them. If the black hole rotates faster than the disk, $ca/GM_H>0.36$ for a thin Keplerian disk, then energy and angular momentum are extracted from the black hole and transferred to the disk ($M_H$ is the mass and $a M_H$ is the angular momentum of the black hole).  This way the energy originating in the black hole may be radiated away by the disk.  The total amount of energy that can be extracted from the black hole spun down from $ca/GM_H = 0.998$ to $ca/GM_H = 0.36$ by a thin Keplerian disk is $\\approx 0.15 M_Hc^2$.  This is larger than $\\approx 0.09 M_Hc^2$ which can be extracted by the Blandford-Znajek mechanism. ",
        "introduction": "Extraction of energy from a black hole or an accretion disk through magnetic braking has been investigated by many people. As a rotating black hole is threaded by magnetic field lines which connect with remote astrophysical loads, energy and angular momentum are extracted from the black hole and transported to the remote loads via Poynting flux (Blandford \\& Znajek 1977; Macdonald \\& Thorne 1982; Phinney 1983). This is usually called the Blandford-Znajek mechanism and has been suggested to be a plausible process for powering jets in active galactic nuclei (Rees, Begelman, Blandford, \\& Phinney 1982; Begelman, Blandford, \\& Rees 1984) and gamma ray bursts (Paczy\\'nski 1993; Lee, Wijers, \\& Brown 1999). Similar process can happen to an accretion disk when some of magnetic field lines threading the disk are open and connect with remote astrophysical loads (Blandford 1976; Blandford \\& Znajek 1977; Macdonald \\& Thorne 1982; Livio, Ogilvie, \\& Pringle 1999; Li 1999).  In this paper we investigate the effects of magnetic field lines connecting a Kerr black hole with a disk surrounding it.  This kind of magnetic field lines are expected to exist and have important effects (Macdonald \\& Thorne 1982; Blandford 1999, 2000; Gruzinov 1999).  We find that, with the existence of such magnetic coupling between the black hole and the disk, energy and angular momentum are transfered between them. If the black hole rotates faster than the disk, energy and angular momentum are extracted from the black hole and transferred to the disk via Poynting flux. This is the case when $a/M_H>0.36$ for a thin Keplerian  disk, where $M_H$ is the mass of the black hole and $aM_H$ is the angular momentum of the black hole. Throughout the paper we use the geometric units with $G = c = 1$.  The energy deposited into the disk by the black hole is eventually radiated to infinity by the disk. This provides a way for extracting energy from a black hole through its disk. If the disk has no accretion (or the accretion rate is very low), the power of the disk essentially comes from the rotational energy of the black hole. We will show that the magnetic coupling between the black hole and the disk has a higher efficiency in extracting energy from a Kerr black hole than the Blandford-Znajek mechanism. ",
        "conclusions": "When a black hole rotates faster than the disk, which is the case if $a/M_H>0.36$ for a Kerr black hole with a thin Keplerian disk, then the black hole exerts a torque at the inner edge of the disk. The torque transfers energy and angular momentum from the black hole to the disk. This is similar to the ``propeller'' mechanism in the case of a magnetized neutron star with a disk (Illarionov \\& Sunyaev 1975). The energy transfered to the disk is eventually radiated to infinity by the disk. This provides a mechanism for extracting energy from a black hole through its disk. For a Kerr black hole with the initial mass $M_H$ and spin $a/M_H = 0.998$, the total amount of energy that can be extracted by a thin Keplerian disk is $\\approx 0.15 M_H$. Therefore, this is more efficient than the Blandford-Znajek mechanism which can extract only $\\approx 0.09 M_H$.  When the black hole rotates slower than the disk, i.e. $0\\le a/M_H<0.36$, energy and angular momentum are transferred from the disk to the black hole, and the disk accretes onto the black hole."
    },
    "0002/astro-ph0002232_arXiv.txt": {
        "abstract": "The structure of neutron stars is considered from theoretical and observational perspectives.  We demonstrate an important aspect of neutron star structure: the neutron star radius is primarily determined by the behavior of the pressure of matter in the vicinity of nuclear matter equilibrium density.  In the event that extreme softening does not occur at these densities, the radius is virtually independent of the mass and is determined by the magnitude of the pressure.  For equations of state with extreme softening, or those that are self-bound, the radius is more sensitive to the mass.  Our results show that in the absence of extreme softening, a measurement of the radius of a neutron star more accurate than about 1 km will usefully constrain the equation of state.  We also show that the pressure near nuclear matter density is primarily a function of the density dependence of the nuclear symmetry energy, while the nuclear incompressibility and skewness parameters play secondary roles.  In addition, we show that the moment of inertia and the binding energy of neutron stars, for a large class of equations of state, are nearly universal functions of the star's compactness.  These features can be understood by considering two analytic, yet realistic, solutions of Einstein's equations, due, respectively, to Buchdahl and Tolman.  We deduce useful approximations for the fraction of the moment of inertia residing in the crust, which is a function of the stellar compactness and, in addition, the presssure at the core-crust interface. \\\\   \\ni {\\em Subject headings:} structure of stars -- equation of state -- stars: interiors -- stars: neutron ",
        "introduction": "The theoretical study of the structure of neutron stars is crucial if new observations of masses and radii are to lead to effective constraints on the equation of state (EOS) of dense matter.  This study becomes ever more important as laboratory studies may be on the verge of yielding evidence about the composition and stiffness of matter beyond the nuclear equilibrium density $\\rho_s\\cong2.7\\cdot10^{14}$ g cm$^{-3}$.  Rhoades \\& Ruffini (1974) demonstrated that the assumption of causality beyond a fiducial density $\\rho_f$ sets an upper limit to the maximum mass of a neutron star: $4.2\\sqrt{\\rho_s/\\rho_f}{\\rm~M}_\\odot$.  However, theoretical studies of dense matter have considerable uncertainty in the high-density behavior of the EOS largely because of the poorly constrained many-body interactions.  These uncertainties have prevented a firm prediction of the maximum mass of a beta-stable neutron star.  To date, several accurate mass determinations of neutron stars are available from radio binary pulsars (Thorsett \\& Chakrabarty 1998), and they all lie in a narrow range ($1.25-1.44$ M$_\\odot$).  One neutron star in an X-ray binary, Cyg X-2, has an estimated mass of $1.8\\pm0.2$ M$_\\odot$ (Orosz \\& Kuulkers 1999), but this determination is not as clean as for a radio binary.  Another X-ray binary, Vela X-1, has a reported mass around 1.9 M$_\\odot$ (van Kerkwijk et al. 1995a), although Stickland et al. (1997) argue it to be about 1.4 M$_\\odot$.  A third object, the eclipsing X-ray binary 4U 1700-37, apparently contains an object with a mass of $1.8\\pm0.4$ M$_\\odot$ (Heap \\& Corcoran 1992), but Brown, Weingartner \\& Wijers (1996) have argued that since this source does not pulse and has a relatively hard X-ray spectrum, it may contain a low-mass black hole instead.  It would not be surprising if neutron stars in X-ray binaries had larger masses than those in radio binaries, since the latter have presumably accreted relatively little mass since their formation. Alternatively, Cyg X-2 could be the first of a new and rarer population of neutron stars formed with high masses which could originate from more massive, rarer, supernovae.  If the high masses for Cyg X-2 or Vela X-1 are confirmed, significant constraints on the equation of state would be realized.  On the other hand, there is a practical, albeit theoretical, lower mass limit for neutron stars, about $1.1-1.2$ M$_\\odot$, which follows from the minimum mass of a protoneutron star.  This is estimated by examining a lepton-rich configuration with a low-entropy inner core of $\\sim0.6$ M$_\\odot$ and a high-entropy envelope (Goussard, Haensel \\& Zdunik 1998).  This argument is in general agreement with the theoretical result of supernova calculations, in which the inner homologous collapsing core material comprises at least 1 M$_\\odot$.  Although accurate masses of several neutron stars are available, a precise measurement of the radius does not yet exist.  Lattimer et al. (1990) (see also Glendenning 1992) have shown that the causality constraint can be used to set a lower limit to the radius: $R\\simge3.04GM/Rc^2$.  For a 1.4 M$_\\odot$ star, this is about 4.5 km.  Estimates of neutron star radii from observations have given a wide range of results.  Perhaps the most reliable estimates stem from observations of thermal emissions from neutron star surfaces, which yield values of the so-called ``radiation radius'' \\begin{equation} R_\\infty=R/\\sqrt{1-2GM/Rc^2}\\,, \\label{rinfty} \\end{equation} a quantity resulting from redshifting the stars luminosity and temperature.  A given value of $R_\\infty$ implies that $R<R_\\infty$ and $M<0.13(R_\\infty/{\\rm~km})$ M$_\\odot$.  Thus, a 1.4 M$_\\odot$ neutron star requires $R_\\infty>10.75$ km.   Those pulsars with at least some suspected thermal radiation generically yield effective values of $R_\\infty$ so small that it is believed that a significant part of the radiation originates from polar hot spots rather than from the surface as a whole.  For example, Golden \\& Shearer (1999) found that upper limits to the unpulsed emission from Geminga, coupled with a parallactic distance of 160 pc, yielded values of $R_\\infty\\simle9.5$ km for a blackbody source and $R_\\infty\\simle10$ km for a magnetized H atmosphere.  Similarly, Schulz (1999) estimated emission radii of less than 5 km, assuming a blackbody for eight low mass X-ray binaries.  Other attempts to deduce a radius include analyses (Titarchuk 1994) of X-ray bursts from sources 4U 1705-44 and 4U 1820-30 which implied rather small values, $9.5<R_\\infty<14$ km.  Recently, Rutledge et al. (1999) found that thermal emission from neutron stars of a canonical 10 km radius was indicated by the interburst emission. However, the modeling of the photospheric expansion and touchdown on the neutron star surface requires a model dependent relationship between the color and effective temperatures.  Absorption lines in X-ray spectra have also been investigated (Inoue 1992) with a view to deducing the neutron star radius.  Candidates for the matter producing the absorption lines are either the accreted matter from the companion star or the products of nuclear burning in the bursts.  In the former case, the most plausible element is thought to be Fe, in which case the relation $R\\approx3.2GM/c^2$, only slightly larger than the minimum possible value based upon causality (Lattimer et al. 1990; Glendenning 1992) is inferred.  In the latter case, plausible candidates are Ti and Cr, and larger values of the radius would be obtained.  In both cases, serious difficulties remain in interpreting the large line widths, of order 100--500 eV, in the $4.1 \\pm 0.1$ keV line observed from many sources.  A first attempt at using light curves and pulse fractions from pulsars to explore the $M-R$ relation suggested relatively large radii, of order 15 km (Page 1995).  However, this method, which assumed dipolar magnetic fields, was unable to satisfactorily reconcile the calculated magnitudes of the pulse fractions and the shapes of the light curves with observations.  The discovery of Quasi-Periodic Oscillations (QPOs) from X-ray emitting neutron stars in binaries provides a possible way of limiting neutron star masses and radii.  These oscillations are manifested as quasi-periodic X-ray emissions, with frequencies ranging from tens to over 1200 Hz.  Some QPOs show multiple frequencies, in particular, two frequencies $\\nu_1$ and $\\nu_2$ at several hundred Hz.  These frequencies are not constant, but tend to both increase with time until the signal ultimately weakens and disappears.  In the beat frequency model (Alpar \\& Shaham 1985, Psaltis et al. 1998), the highest frequency $\\nu_2$ is associated with the Keplerian frequency $\\nu_K$ of inhomogeneities or blobs in an accretion disc.  The largest such frequency measured to date is $\\nu_{max}=1230$ Hz.  However, general relativity predicts the existence of a maximum orbital frequency, since the inner edge of an accretion disc must remain outside of the innermost stable circular orbit, located at a radius of $r_s=6GM/c^2$ in the absence of rotation.  This corresponds to a Keplerian orbital frequency of $\\nu_s=\\sqrt{GM/r_s^3}/2\\pi$ if the star is non-rotating.  Equating $\\nu_{max}$ with $\\nu_s$, and since $R<r_s$, one deduces \\begin{equation} M\\simle1.78{\\rm~M}_\\odot;\\qquad R\\simle8.86 (M/{\\rm M}_\\odot){\\rm~km}.\\label{qpo} \\end{equation} Corrections due to stellar rotation are straightforward to deduce and produce small changes in these limits (Psaltis et al. 1998).  The lower frequency $\\nu_1$ is associated with a beat frequency between $\\nu_2$ and the spin frequency of the star.  This spin frequency is large enough, of order 250-350 Hz, to alter the metric from a Schwarzschild geometry, and increases the theoretical mass limit in equation (\\ref{qpo}) to about 2.1 M$_\\odot$ (Psaltis et al. 1998). This remains strictly an upper limit, however, unless further observations support the interpretation that $\\nu_{max}$ is associated with orbits at precisely the innermost stable orbit.  If the frequency $\\nu_2-\\nu_1$ is to be associated with the spin of the neutron star, it should remain constant in time.  However, recent observations reveal that it changes with time in a given source. Osherovich \\& Titarchuk (1999) proposed a model in which $\\nu_1$ is the Keplerian frequency and $\\nu_2$ is a hybrid frequency of the Keplerian oscillator under the influence of a magnetospheric Coriolis force.  In this model, the frequencies are related to the neutron star spin frequency $\\nu$ by \\begin{equation} \\nu_2=\\sqrt{\\nu_1^2+(\\nu/2\\pi)^2}. \\label{ko} \\end{equation} Osherovich \\& Titarchuk argue that this relation fits the observed variations of $\\nu_2$ and $\\nu_1$ in several QPOs.  The Keplerian frequency in Osherovich \\& Titarchuk's model, being associated with the lower frequency $\\nu_1$, however, is at most 800 Hz, leading to an upper mass limit that is nearly 3 M$_\\odot$ and is therefore of little practical use to limit either the star's mass or radius.  An alternative model, proposed by Stella \\& Vietri (1999), associates $\\nu_2$ with the Keplerian frequency of the inner edge of the disc, $\\nu_K$, and $\\nu_2-\\nu_1$ with the precession frequency of the periastron of slightly eccentric orbiting blobs at radius $r$ in the accretion disc.  In a Schwarzschild geometry, $\\nu_1=\\nu_K\\sqrt{1-6GM/rc^2}$.  Note that $(\\nu_K-\\nu_1)^{-1}$ is the timescale that an orbiting blob recovers its original orientation relative to the neutron star and the Earth, so that variations in flux are expected to be observed at both frequencies $\\nu_K$ and $\\nu_K-\\nu_1$.  Presumably, even eccentricities of order $10^{-4}$ lead to observable effects.  This model predicts that \\begin{equation} \\nu_1/\\nu_2=1-\\sqrt{1-6(GM\\nu_2)^{2/3}/c^2}, \\label{sv} \\end{equation} a relation that depends only upon $M$.  Equation~(\\ref{sv}) can also fit observations of QPOs, but only if $1.9\\simle M/{\\rm M}_\\odot\\simle 2.1$.  This result is not very sensitive to complicating effects due to stellar rotation: the Lense-Thirring effect and oblateness.  This mechanism only depends on gravitometric effects, and may apply also to accreting black hole systems (Stella, Vietri \\& Morsink 1999).  Prospects for a radius determination have improved in recent years with the discovery of a class of isolated, non-pulsing, neutron stars. The first of these is the nearby object RX J185635-3754, initially discovered in X-rays (Walter, Wolk \\& Neuha\\\"user 1996) and confirmed with the Hubble Space Telescope (Walter \\& Matthews 1997). The observed X-rays, from the ROSAT satellite, are consistent with blackbody emission with an effective temperature of about 57 eV and relatively little extinction.  The fortuitous location of the star, in the foreground of the R CrA molecular cloud, coupled with the small levels of extinction, limits the distance to $D<120$ pc.  The fact that the source is not observable in radio and its lack of variability in X-rays implies that it is not a pulsar, unlike other identified radio-silent isolated neutron stars.  This gives the hope that the observed radiation is not contaminated with non-thermal emission as in the case for pulsars.  The X-ray flux of RX~J185635-3754, combined with a best-fit blackbody $T_{eff}=57$ eV, yields $R_\\infty\\approx 7.3 (D/120 {\\rm~pc}){\\rm~ km}$.  Such a value for $R_\\infty$, even coupled with the maximum distance of 120 pc, is too small to be consistent with any neutron star with more than 1 M$_\\odot$.  But the optical flux is about a factor of 4 brighter than what is predicted by the X-ray blackbody. The reconciliation the X-ray and optical fluxes through atmosphere modeling naively implies an increase in $R_\\infty$ of approximately $4^{2/3}\\cong2.5$.  (This results since the X-ray flux is proportional to $R_\\infty^2T_{eff}^4$, while the optical flux is on the Rayleigh-Jeans tail of the spectrum and is hence proportional to $R_\\infty^2T_{eff}$.  One seeks to enhance the predicted optical flux by 4 while keeping the X-ray flux fixed, as this is approximately equal to the total flux.) An et al. (2000) determined for non-magnetized heavy element atmospheres that $R_\\infty/D\\cong 0.18\\pm0.05$ km pc$^{-1}$, which is rough agreement with the above naive expectations.  However, uncertainties in the atmospheric composition and the quality of the existing data precluded obtaining a more precise estimate of $R_\\infty/D$.  An et al. concluded, in agreement with expectations based upon the general results of Romani (1987) and Rajagopal, Romani \\& Miller (1997), that the predicted spectrum of a heavy element atmosphere, but not a light element atmosphere, was consistent with all the observations.  This is in contrast to the conclusions of Pavlov et al. (1996), whose results for RX J185635-3754 implied that the observations in the optical and X-ray bands were incompatible with atmospheric modelling for both heavy element and light element non-magnetized atmospheres, unless the distance to this star is greater than the presumed maximum of 120 pc based upon the star's location in front of the R Cor Aus molecular cloud.  Future prospects for determining the radius of this neutron star are discussed in \\S 7.  Our objectives in this paper are 1) to demonstrate specifically how the accurate measurement of a neutron star radius would constrain the dense matter EOS, and 2) to provide general relationships for other structural quantities, such as the moment of inertia and the binding energy, that are relatively EOS-independent, and which could be used to constrain the neutron star mass and/or radius.  We will examine a wide class of equations of state, including those that have extreme softening at high densities.  In addition, we will examine analytic solutions to Einstein's equations which shed light on the results we deduce empirically.  In all cases, we will focus on non-rotating, non-magnetized neutron stars at zero temperature.  Lindblom (1992) had suggested that a series of mass and radius measurements would be necessary to accurately constrain the dense matter equation of state.  His technique utilizes a numerical inversion of the neutron star structure equation.  Our results instead suggest that important constraints on the EOS can be achieved with even a single radius measurement, if it is accurate enough, and that the quality of the constraint is not very sensitive to the mass.  The fact that the range of accurately determined neutron star masses is so small, only about 0.2 M$_\\odot$ to date, further implies that important constraints can be deduced without simulaneous mass-radius measurements.  Of course, several measurements of neutron star masses and radii would greatly enhance the constraint on the equation of state.   In \\S~2, the equations of state selected in this paper are discussed. In \\S~3, the mass-radius relation for a sample of these equations of state are discussed.  A quantitative relationship between the radii of normal neutron stars and the pressure of matter in the vicinity of $n_s$ is empirically established and theoretically justified.  In turn, how the matter's pressure at these densities depends upon fundamental nuclear parameters is developed.  In \\S~4, analytic solutions to the general relativistic equations of hydrostatic equilibrium are explored.  These lead to useful approximations for neutron star structure and which directly correlate other observables such as moments of inertia and binding energy to the mass and radius. It is believed that the distribution of the moment of inertia inside the star is crucial in the interpretation of glitches observed in the spin down of pulsars, so that measurements of the sizes and frequencies of glitches can constrain neutron star masses and radii (Link, Epstein \\& Lattimer 1999).  In \\S~5, expressions for the fraction of moment of inertia contained within the stellar crust, as a function of mass, radius, and equation of state, are derived.  In \\S~6, expressions for the binding energy are derived.  \\S~7 contains a summary and outlook. ",
        "conclusions": "We have demonstrated the existence of a strong correlation between the pressure near nuclear saturation density inside a neutron star and the radius which is relatively insensitive to the neutron star's mass and equation of state for normal neutron stars.  In turn, the pressure near the saturation density is primarily determined by the isospin properties of the nucleon-nucleon interaction, specifically, as reflected in the density dependence of the symmetry energy, $S_v(n)$. This result is not sensitive to the other nuclear parameters such as $K_s$, the nuclear incompressibility parameter, or $K^\\prime_s$, the skewness parameter.  This is important, because the value of the symmetry energy at nuclear saturation density and the density dependence of the symmetry energy are both difficult to determine in the laboratory.  Thus, a measurement of a neutron star's radius would yield important information about these quantities.  Any measurement of a radius will have some intrinsic uncertainty.  In addition, the empirical relation we have determined between the pressure and radius has a small uncertainty.  It is useful to display how accurately the equation of state might be established from an eventual radius measurement.  This can be done by inverting equation (\\ref{correl}), which yields \\begin{equation} P(n) \\simeq [R_M/C(n,M)]^4\\,. \\label{correli} \\end{equation} The inferred ranges of pressures, as a function of density and for three possible values of $R_{1.4}$, are shown in Figure~\\ref{fig:err}. It is assumed that the mass is 1.4 M$_\\odot$, but the results are relatively insensitive to the actual mass.  Note from Table 3 that the differences between $C$ for 1 and 1.4 M$_\\odot$ are typically less than the errors in $C$ itself.  The light shaded areas show the pressures including only errors associated with $C$.  The dark shaded areas show the pressures when a hypothetical observational error of 0.5 km is also taken into account.  These results suggest that a useful restriction to the equation of state is possible if the radius of a neutron star can be measured to an accuracy better than about 1 km.  The reason useful constraints might be obtained from just a single measurement of a neutron star radius, rather than requiring a series of simultaneous mass-radii measurements as Lindblom (1992) proposed, stems from the fact that we have been able to establish the empirical correlation, equation (\\ref{correl}).  In turn, it appears that this correlation exists because most equations of state have slopes $d \\ln P/d \\ln n\\simeq2$ near $n_s$.  The best prospect for measuring a neutron star's radius may be the nearby object RX J185635-3754.  It is anticipated that parallax information for this object will be soon available (Walter, private communication).  In addition, it may be possible to identify spectral lines with the Chandra and XMM X-ray facilities that would not only yield the gravitational redshift, but would identify the atmospheric composition.  Not only would this additional information reduce the uncertainty in value of $R_\\infty$, but, {\\em both} the mass and radius for this object might thereby be determined.  It is also possible that an estimate of the surface gravity of the star can be found from further comparisons of observations with atmospheric modelling, and this would provide a further check on the mass and radius.  We have presented simple expressions for the moment of inertia, the binding energy, and the crustal fraction of the moment of inertia for normal neutron stars which are largely independent of the EOS. If the magnitudes of observed glitches from Vela are connected with the crustal fraction of moment of inertia, the formula we derived establishes a more stringent limit on the radius than causality.   We thank A. Akmal, V. R. Pandharipande and J. Schaffner-Bielich for making the results of their equation of state calculations available to us in tabular form.  This work was supported in part by the USDOE grants DOE/DE-FG02-87ER-40317 \\& DOE/DE-FG02-88ER-40388."
    },
    "0002/hep-ph0002257_arXiv.txt": {
        "abstract": "New physics beyond the electroweak scale may increase weak interaction cross sections beyond the Standard Model predictions.  Such cross sections can be expected within theories that solve the hierarchy problem of known interactions with a unification scale in the TeV range. We derive constraints on these cross sections from the flux of neutrinos expected from cosmic ray interactions with the microwave background and the non-observation of horizontal air showers. We also discuss how this limit can be improved by upcoming cosmic ray and neutrino experiments, and how the energy dependence of the new interactions can be probed by these experiments. ",
        "introduction": "It has been suggested that the neutrino-nucleon cross section, $\\sigma_{\\nu N}$, can be enhanced by new physics beyond the electroweak scale in the center of mass (CM) frame, or above about a PeV in the nucleon rest frame. At ultra-high energies (UHE),  neutrinos can in principle acquire $\\nu N$ cross sections approaching hadronic levels.  It may therefore be possible for UHE neutrinos to initiate air showers, which would offer a very direct means of probing these new UHE interactions.  The results of this kind of experiment have important implications for both neutrino astronomy and high energy physics.  For the lowest partial wave contribution to the cross section of a point-like particle, this new physics would violate unitarity~\\cite{bhg}. However, two major possibilities have been discussed in the literature for which unitarity bounds need not be violated. In the first, a broken SU(3) gauge symmetry dual to the unbroken SU(3) color gauge group of strong interaction is introduced as the ``generation symmetry'' such that the three generations of leptons and quarks represent the quantum numbers of this generation symmetry. In this scheme, neutrinos can have close to strong interaction cross sections with quarks. In addition, neutrinos can interact coherently with all partons in the nucleon, resulting in an effective cross section comparable to the geometrical nucleon cross section.  This model lends itself to experimental verification through shower development altitude statistics, as described by its authors~\\cite{bhfpt}.  The present paper will not affect the plausibility of this possibility, which largely awaits more UHE cosmic ray (UHECR) events to compare against.  The second possibility consists of a large increase in the number of degrees of freedom above the electroweak scale~\\cite{kovesi-domokos}. A specific implementation of this idea is given in theories with $n$ additional large compact dimensions and a quantum gravity scale $M_{4+n}\\sim\\,$TeV that has recently received much attention in the literature~\\cite{tev-qg} because it provides an alternative solution (i.e., without supersymmetry) to the hierarchy problem in grand unifications of gauge interactions. For parameters consistent with measured cross sections at electroweak energies and below, at $\\simeq10^{20}\\,$eV, the $\\nu N$ cross section can approach $10^{-4}-10^{-2}$ times a typical strong cross section~\\cite{jain}. Therefore, in these models the mean free path for $\\nu N$ interactions at these energies is on the order of kilometers to hundreds of kilometers in air. Experimentally, this would be reflected by a specific energy dependence of the typical column depth at which air showers initiate. As we will show, this leads to a useful signature, deriving from zenith angle variations with shower energy.  Comparison with observations would either support these scenarios or constrain them in a way complementary to studies in terrestrial accelerators~\\cite{coll-extra-dim} or other laboratory experiments that have recently appeared in the literature.  Both scenarios for increasing $\\nu N$ cross sections raise the question of how far UHE neutrinos can travel in space.  In the first case, due to the {\\it coherent} nuclear interactions, the cross section approaches the nucleon-nucleon cross section, and so intergalactic matter becomes the greatest threat to the neutrino traveling long distances unimpeded.  However, an order of magnitude calculation shows that the mean free path for such a neutrino is $\\sim10^2$~Gpc, if we simply assume the matter density of the Universe is its critical density.  Within a galaxy with similar interstellar density to the Milky Way, the mean free path is $\\sim\\,$Mpc; the UHE neutrino is not attenuated before reaching Earth's atmosphere. For the extra dimension model, the cross sections are lower, so interaction with matter is not a problem.  With Cosmic Microwave Background (CMB) photons, the CM energy is too low to give rise to a significant interaction probability. If neutrinos are massive, then the relic neutrino background is more of a threat.  But assuming the neutrino mass is below 92 eV, as conservatively required to prevent over-closing the Universe for a stable species~\\cite{kolb}, the mean free path is $\\sim10^4$~Gpc at a CM energy of $\\sim100\\,$GeV, corresponding to a $\\sim10^{20}\\,$eV neutrino. The cosmological horizon is only $\\sim15$~Gpc, so for both scenarios the neutrinos can reach Earth unhindered from arbitrarily great distances.  A major motivation for attempts to augment the $\\nu N$ cross section at ultra-high energies derives from the paradox of UHECRs. Above $\\sim10^{20}$~eV, protons and neutrons see CMB photons doppler shifted to high energy gamma rays, and encounters generate pions, thereby reducing the resulting nucleon's energy by at least the pion rest mass.  One therefore expects the Greisen-Zatespin-Kuzmin (GZK) cutoff~\\cite{gzk}, wherein cosmic rays above $E_{\\rm GZK}\\simeq7\\times10^{19}$~eV would lose their energy within $\\simeq50\\,$Mpc. However, since we observe cosmic rays at super-GZK energies, it has been proposed that these UHECRs might really be neutrinos rather than nucleons (for example, see~\\cite{bhfpt}).  While this is a strong motivation to explore new UHE neutrino interactions, we will not need to assume that UHECRs are neutrinos.  The paper is outlined as follows: In Section~II we derive a constraint on the neutrino-nucleon cross section resulting from the non-observation of horizontal air showers induced by neutrinos produced by interactions of the highest energy cosmic rays with the cosmic microwave background. In Section~III we discuss potential improvements of this constraint including other hypothesized neutrino sources and next generation experiments. In Section~IV we narrow the discussion to the case of a $\\nu N$ cross section scaling linearly with CM energy, as an example. A zenith angle profile is presented as a signature of this scaling, and a means of constraining it. We conclude in Section~V. ",
        "conclusions": "A direct measurement approach has been presented here for bounding the extent to which UHE neutrinos acquire greater interactions with matter than the Standard Model prescribes.  By looking for nearly horizontal air showers, we can limit the UHE neutrino flux striking Earth.  When this is combined with a natural source of cosmogenic UHE neutrinos, a model-independent upper limit on the cross section results. Air shower experiments can also test for the energy dependence of cross sections in the range $\\sim10^{-29}-10^{-27}\\,{\\rm cm}^2$; we have discussed the specific case of a linear energy dependence, as motivated by some large extra compact dimensions models that have enjoyed much recent attention in the community. In such scenarios, fundamental quantum gravity scales in the $1-10\\,$TeV range can be probed.  This could also be relevant for testing string theory scenarios with string scales in the TeV range, however reliable theoretical predictions for the relevant cross sections in such scenarios are not yet available.  If this method proves beneficial, knowledge will be gained about new interactions at ultra-high energies. Should the data support extra dimensions, our understanding of UHE interactions between {\\it all} particles (not just $\\nu N$) will be changed because changes in gravity affect every particle. In general, cosmic neutrinos and UHECRs will probe interactions at energies a few orders of magnitude beyond what can be achieved in terrestrial accelerators in the forseeable future."
    },
    "0002/astro-ph0002248_arXiv.txt": {
        "abstract": "We update the Big Bang Nucleosynthesis calculations on the basis of the recent NACRE compilation of reaction rates. The average values of the calculated abundances of light nuclei do not differ significantly from those obtained using the previous Fowler's compilation. However,  ${^7}Li$ is slightly depressed at high baryon to photon ratio $\\eta$. Concerning ${^{10}}B$, its abundance is significantly lower than the one calculated with the Caughlan and Fowler (1988) rates as anticipated by Rauscher and Raimann (1997). We estimate the uncertainties related to the nuclear reaction rates on the abundances of $D$, ${^3}He$, ${^4}He$, ${^6}Li$, ${^7}Li$, ${^9}Be$, ${^{10}}B$ and ${^{11}}B$ of cosmological and astrophysical interest. The main uncertainty concerns the $D(p,\\gamma){^3}He$  reaction rate affecting the synthesis of ${^7}Li$ at rather high baryonic density and also the ${^3}He(\\alpha,\\gamma){^7}Be$ and ${^7}Li(p,\\alpha){^4}He$ reactions. On the left part of the lithium valley the uncertainty is reduced due to the improvement of the measurement of the $T(\\alpha,\\gamma)^{7}Li$ reaction rate. The observed abundances of the nuclei of interest are compared to the predictions of the BBN model, taking into account both observational and theoretical uncertainties. Indeed, the ${^7}Li$ abundance observed in halo stars (Spite plateau) is now  determined with high precision since the thickness of this plateau appears, in the light of recent observations, exceptionnaly small ($<$ 0.05 dex). The potential destruction/dilution of ${^7}Li$ in the outer layers of halo stars which could mask the true value of the primordial abundance is in full debate, but the present trend is towards a drastic reduction of the depletion factor (about 0.10 dex). It is why we use this isotope as a preferred baryometer. Even though much efforts have been devoted to the determination of deuterium in absorbing clouds in the line of sight of remote quasars, the statistics is very poor compared to the long series of lithium measurements. Taking into account these lithium constraints, two possible baryonic density ranges emerge, $\\eta_{10}= 1.5 - 1.9$ and $\\eta_{10} = 3.3 - 5.1$. In the first case, $Li$ is in concordance with $D$ from Webb et al (1997) and ${^4}He$ from Fields and Olive (1998) and Peimbert and Peimbert (2000). In the second case, agreement is achieved with $D$ from Tytler et al (2000) and ${^4}He$ from Izotov and Thuan (1998).  Concerning the less abundant light isotopes, the theoretical BBN abundance of ${^6}Li$ is affected by a large uncertainty due to the poor knowledge of the $D(\\alpha,\\gamma){^6}Li$ reaction rate. However, at high $\\eta$, its abundance is so low that there is little chance to determine observationally the true BBN ${^6}Li$ abundance. But, at low $\\eta$, its abundance being one thousandth of that of primordial ${^7}Li$, 6/7 ratio measurements at very low metallicity are not totally hopeless in the future. Nevertheless, in the present situation, ${^6}Li$ is cosmologically relevant, though indirectly, since its mere presence in a few halo stars, corroborates the fact that it is essentially intact in these stars together with $^{7}Li$ and thus the Spite plateau can be used as such to infer the primordial ${^7}Li$ abundance. The $Be$ and $B$ abundances produced in the Big Bang  are orders of magnitudes lower, and spallation of fast carbon and oxygen is probably their unique source, in the early Galaxy. ",
        "introduction": "Besides the expansion of the Universe and the ubiquitous presence of the fossil cosmological radiation, the Big Bang Nucleosynthesis (BBN) is one pilar of modern cosmology. It allows, in principle, the derivation of the baryonic density of the universe (see for reviews Schramm and Turner 1998, Sarkar 1996, 1999 and Olive et al 2000). The determination of light element abundances has improved dramatically in the recent past and the planned observations of $D$, ${^4}He$, ${^6}Li$ and ${^7}Li$ should allow a precise determination (10\\% is a reachable goal) of the universal baryonic density, provided the precision of the model is made compatible with this objective. Taking advantage of the  release of a new compilation of thermonuclear reation rates, called NACRE (Nuclear Astrophysics Compilation of REaction rates) (Angulo et al 1999), we have updated the standard BBN model developed at the Institut d'Astrophysique de Paris, including the analysis of the $Be$ and $B$ production in the Big Bang.  On the other hand, observations of light isotopes have florished: i) $D/H$ has been observed in absorbing clouds on the sightline of remote quasars, ii) refined observations of $^{4}He$ have been performed in extragalactic very metal poor HII regions and in the Small Magellanic Cloud (Peimbert and Peimbert 2000), iii) the $^{6}Li$ abundance has been determined in two halo stars, iv) high quality $^{7}Li$ observations in the halo stars have been accumulated. A review of the present data can be found in Olive et al (2000) and Tytler et al (2000). Thus, it is timely to reassess the determination of the baryonic density of the Universe in the light of advances in nuclear physics and astronomical observations.  In section II, we present the new compilation of the reaction rates and compare it with the classical Caughlan-Fowler (1988) ones; we evaluate the sensitivity of the different light element abundances to the change of each relevant reaction rate; BBN calculations are performed using i) recommended values of the reaction rates and ii) extreme values obtained from the low and high rate limits. In section III, we discuss the astrophysical status of each isotope, both observationally and theoretically in order to confront it to the BBN calculation; we deduce the baryonic density of the Universe. Finally, we draw conclusion in section IV and stress the importance of precise measurements of a few key nuclear cross sections and refined abundance determinations of $D$ and $^{6}Li$. ",
        "conclusions": "Big bang nucleosynthesis has been studied since a long time, but it deserves permanent care since it gives access to the baryon density which is a key cosmological parameter. This work has been aimed at integrating the last development in both fields of nuclear physics and observational abundance determination of light elements.   1. The update of the reaction rates of the BBN using the NACRE compilation does not lead to crucial modifications of the general conclusions concerning the baryonic content of the Universe. The average values of the abundances of isotopes of cosmological interest are in general similar to that calculated on the basis of the Caughlan-Fowler (1988) compilation.  2. However, the modification of the $D(p,\\gamma)^{3}He$ reaction rate leads to a $^{7}Li$ abundance slightly lower at $\\eta > 3$. But the uncertainty on this reaction rate remains high ($\\pm$30\\%). At $\\eta< 3$, the revision of the $T(\\alpha,\\gamma)^{7}Li$ reaction rate leads to a very neat reduction of the uncertainty of the calculated $^{7}Li$ abundance. However this reaction, together with the $^{7}Li(p,\\alpha)^{4}He$, remain the main sources of the $^{7}Li$ uncertainty at low $\\eta$.  3. The abundance of $^{10}B$ is modified by the new $^{10}B(p,\\alpha)^{7}Li$ reaction rate, but there is no cosmological consequences.  4. $^{6}Li$ is affected by the large uncertainty of the $D(\\alpha,\\gamma)^{6}Li$ reaction. However, $^{6}Li$ is essentially of spallative origin.  5. Owing to the high observational reliability of the $^{7}Li$ abundance data with respect to the $D$ data avalaible both rare and debated, we choose it as the leading baryometer, since it appears that the primitive $^{7}Li$ is almost intact in halo stars. Due to the competition between $T(\\alpha,\\gamma)^{7}Li$ and $D(p,\\gamma)^{3}He(\\alpha,\\gamma)^{7}Be$ (e $\\nu$) $^{7}Li$, the curve of $^{7}Li$ vs $\\eta$ presents a valley shape. Consequently, the observational error box of $^{7}Li$ leads to two ranges of $\\eta$ : $1.5<\\eta<1.9$ and $3.3<\\eta<5.1$ (corresponding to $0.013<\\Omega_B<0.019$ and $0.029<\\Omega_B<0.045$ for h=0.65). In both cases, all the dark matter in the halo of our Galaxy could be baryonic. However, only a fraction of 20\\% is detected  through microlensing events in the direction of the Magellanic clouds.  6. These two $\\eta$ ranges are confronted to the other available cosmologically relevant isotopes, namely $D$ and $^{4}He$. The first $\\eta$ range agrees with a high $D$ and low $^{4}He$ values, the second range is in concordance with a low $D$ and high $^{4}He$ values. At present, none of these solutions can be excluded.  7. In the future, on the nuclear physics front, it would be important to (re-)measure the $D(\\alpha,\\gamma)^{6}Li$ reaction to reduce the uncertainty on the calculated $^{6}Li$ abundance. High precision measurements over the full energy range of interest to BBN  of the $D(p,\\gamma)^3He$, $^3He(\\alpha,\\gamma)^7Be$ and $^7Li(p,\\alpha)^4He$ reactions would also reduce the uncertainty on $^7Li$ abundance calculations. On the astronomical front, more data on $D$ in absorbing clouds on the sightlines of quasars at different redshifts are mandatory to remove the ambiguity. However, if the large scale $D$ dichotomy remains, it will be time to invoke specific mechanisms of $D$ production and destruction like photodisintegration of $D$ and $^{4}He$ by $\\gamma$ ray quasars (blazars) (Cass\\'e and Vangioni-Flam 1997) and/or nuclear spallation. The observations of $^{6}Li$ in two halo stars has been a great progress and a determination of its abundance in very metal poor stars should be pursued. It will help to constrain even more stringently the possible lithium depletion in these stars and confort definitively the primordial status of the Spite plateau. Together with nuclear improvements, refined $^{6}Li$ measurements in very metal poor stars (possibly via the VLT) could perhaps lead us towards the primordial $^{6}Li$ abundance.  We warmly thank Roger Cayrel for illuminating discussions, Jurgen Kiener, Gilles Bogaert and Carmen Angulo for their comments on nuclear data."
    },
    "0002/astro-ph0002281_arXiv.txt": {
        "abstract": "In this paper, Gaussianity of eigenmodes and non-Gaussianity in the Cosmic Microwave Background (CMB) temperature fluctuations in two smallest compact hyperbolic (CH) models are investigated. First, it is numerically found that the expansion coefficients of low-lying eigenmodes on the two CH manifolds behave as if they are Gaussian random numbers at almost all the places. Next, non-Gaussianity of the temperature fluctuations in the ($l,m$) space in these models is studied. Assuming that the initial fluctuations are Gaussian, the real expansion coefficients $b_{l m}$ of the temperature fluctuations in the sky are found to be distinctively non-Gaussian. In particular, the cosmic variances are found to be much larger than that for Gaussian models. On the other hand, the anisotropic structure is vastly erased if one averages the fluctuations at a number of different observing points because of the Gaussian pseudo-randomness of the eigenmodes. Thus the dominant contribution to the two-point correlation functions comes from the isotropic terms described by the angular power spectra $C_l$. Finally, topological quantities: the total length and the genus of isotemperature contours are investigated. The variances of total length and genus at high and low threshold levels are found to be considerably larger than that of Gaussian models while the means almost agree with them. ",
        "introduction": "\\indent In recent years, locally isortopic and homogeneous Friedmann-Robertson-Walker (FRW) models with non-trivial topology have attracted much attention. In the standard scenario, simply-connectivity of the spatial hypersurface is assumed for simplicity.  However, the Einstein equations, being local equations, do not fix the global topology of the spacetime. In other words, a wide variety of topologically distinct spacetimes with the same local geometry described by a local metric element remain unspecified (see \\cite{Lachieze} for review on the cosmological topology). The determination of the global topology of the universe is one of the most important problem of the modern observational cosmology. \\\\ \\indent For flat models without the cosmological constant, severest constraints have been obtained by using the COBE DMR data. The suppression of the fluctuations on scales beyond the topological identification scale $L$ leads to the decrease of the angular power spectra $C_l$ of the Cosmic Microwave Background (CMB) temperature fluctuations on large angular scales which puts a lower bound $L\\!\\ge\\! 2400~ h^{-1} $Mpc (with $h\\!=\\!H_{0}/100~\\mathrm{kms^{-1}Mpc^{-1} }$) for a compact flat 3-torus model without the cosmological constant \\cite{Stevens,Oliveira}. Similar constraints have been obtained for other compact flat models \\cite{Levin}. The maximum expected number of copies of the fundamental domain (cell) inside the last scattering surface is approximately 8 for the 3-torus model. \\\\ \\indent In contrast, for low density models, the constraint could be considerably milder than the locally isotropic and homogeneous flat (Einstein-de-Sitter) models since a bulk of large-angle CMB fluctuations can be produced by the so-called (late) integrated Sachs-Wolfe effect (ISW) \\cite{HSS,Cornish2} which is the gravitational blueshift effect of the free streaming photons by the decay of the gravitational potential. As the gravitational potential decays in either $\\Lambda$-dominant epoch or curvature dominant epoch, the free streaming photons with large wavelength (the light travel time across the wavelength is greater than or comparable to the decay time) that climbed a potential well at the last scattering experience blueshifts due to the contraction of the comoving space along the trajectories of the photons. Because the angular sizes of the fluctuations produced at late time are large, the suppression of the fluctuations on scale larger than the topological identification scale does not lead to a significant suppression of the large-angle power if the ISW effect is dominant. Recent works \\cite{Aurich3,Inoue2,Inoue3,Cornish1} have shown that the large-angle power ($2\\!\\le\\!l\\!\\le\\!20$) are completely consistent with the COBE DMR data for compact hyperbolic (CH) models which include a small CH orbifold, the Weeks and the Thurston manifolds with volume $0.72, 0.94$ and $0.98$ in unit of the cube of the curvature radius, respectively. Note that the Weeks manifold is the smallest and the Thurston manifolds is the second smallest in the known CH manifolds. For instance, the number of copies of the fundamental domain inside the last scattering surface at present is approximately 190 for a Weeks model with $\\Omega_{0}\\!=\\!0.3$. \\\\ \\indent If the space is negatively curved, for a fixed number of the copies of the fundamental domain inside the present horizon, the large-angle fluctuations can be produced much effectively. In negatively curved spaces (hyperbolic spaces), trajectories of photons subtend a much smaller angle in the sky for a given scale. In other words, for a given angle of a pair of two photon trajectories, the physical distance of the trajectories is much greater than that in a flat space. Therefore, even if there is a number of copies of the fundamental domain which intersect the last scattering surface, the number of copies which intersect the wave front (a sphere with $z\\!=\\!\\textrm{const.}$) of the free streaming photons is exponentially decreased at late time when the large-angle fluctuations are produced due to the ISW effect. \\\\ \\indent However, one may not be satisfied with the constraints using only the angular power spectrum $C_{l}$ since it contains only isotropic information of the ensemble averaged temperature fluctuations \\cite{Bond}. If they have anisotropic structures, non-Gaussian signatures must be revealed. In fact, the global isotropy of the locally isotropic and homogeneous FRW models is generally broken. For instance, a flat 3-torus obtained by identifying the opposite faces of a cube is obviously anisotropic at any points. Thus the temperature fluctuations averaged over the initial conditions in these multiply-connected FRW models are no longer $SO(3)$ invariant at a certain point. The temperature fluctuations on the sky are written in terms of (real) spherical harmonics $Q_{lm}(\\n)$ as \\BE \\f{\\Delta T}{T}(\\n)=\\sum_{l}~\\sum_{m=-l}^{l} b_{lm} Q_{lm}(\\n). \\EE If the distribution functions of the real expansion coefficients $b_{lm}$ are $SO(3)$ invariant, the temperature fluctuations must be Gaussian provided that $b_{lm}$'s are independent random numbers \\cite{Mag1}. Therefore, the temperature fluctuations at a certain point in the multiply-connected FRW models are not Gaussian if $b_{lm}$'s are independent. \\\\ \\indent For the simplest flat 3-torus models (without rotations in the identification maps) which are globally homogeneous, it is sufficient to choose one observing point and estimate how the power is distributed among the $m$'s for a given angular scale $l$ in order to see the effect of the global anisotropy. However, in general, one must consider an ensemble of fluctuations at different observing points because of the spatial (global) inhomogeneity. Previous analyses have not fully investigated the dependence of the temperature fluctuations on choice of the observing points. \\\\ \\indent Lack of analytical results on the eigenmodes makes it difficult to investigate the nature of the temperature fluctuations in CH  models. However, we may expect a high degree of complexity in the eigenmodes since the corresponding classical systems (geodesic flows) are strongly chaotic. In fact, it has been numerically found that the expansion coefficients of the low-lying eigenmodes on the Thurston manifold at the point where the injectivity radius is maximal are Gaussian pseudo-random numbers \\cite{Inoue1} which supports the previous analysis of the excited states (higher modes) of a two-dimensional asymmetrical CH model \\cite{Aurich2}. We have put a prefix ''pseudo'' since the eigenmodes are actually constrained by the periodic boundary conditions. These results imply that the statistical properties of the eigenmodes on CH spaces (orbifolds and manifolds) can be described by random-matrix theory (RMT)\\cite{Meh,Boh}. An investigation of the dependence of the property on the observing points is also important since CH spaces have symmetries (isometric groups) which may veil the random feature of the eigenmodes. In this paper, a detailed analysis on the statistical property of low-lying eigenmodes on the Weeks and the Thurston manifolds is conducted. \\\\ \\indent Assuming that the eigenmodes are Gaussian, one can expect that the anisotropic structure in the $(l,m)$ space is vastly erased when one averages the fluctuations over the space. This seems to be a paradox since the CH spaces are actually globally anisotropic. However, one should consider a spatial average of fluctuations with different initial conditions if one believes the Copernican principle that we are not in the center of the universe. Even if the space is anisotropic at a certain point, the averaged fluctuations may look isotropic by considering an ensemble of fluctuations at all the possible observing points. Note that the eigenmodes on CH spaces have no particular directions if they are Gaussian. \\\\ \\indent If the initial fluctuations are constant for each eigenmode, as we shall see, the Gaussian randomness of the temperature fluctuations can be solely attributed to the Gaussian pseudo-randomness of the eigenmodes. In this case, the Gaussian randomness of the temperature fluctuations has its origin in the geometrical property of the space (\\textit{Geometric Gaussianity}). Choosing an observing point is equivalent to fixing a certain initial condition. However, it is much natural to assume that the initial fluctuations are also random Gaussian as the standard inflationary scenarios predict. Then the temperature fluctuations may not obey the Gaussian statistics because they are written in terms of products of two different independent Gaussian numbers rather than sums while they remain almost spatially isotropic if averaged over the space. \\\\ \\indent In this paper, Gaussianity of eigenmodes and non-Gaussianity in the CMB for two smallest CH models (the Weeks and the Thurston models) are investigated. In Sec.~II, numerical results on Gaussianity of eigenmodes are shown and we discuss to what extent the results are generic. In Sec.~III, we study the non-Gaussian behavior of the temperature fluctuations in the ($l,m$) space. In Sec.~IV, topological quantities (total length and genus) of isotemperature contours are numerically simulated for studying the non-Gaussian behavior in the real space. Finally, we summarize our conclusions in Sec.~V. ",
        "conclusions": "In this paper, Gaussianity of the eigenmodes and non-Gaussianity in the CMB temperature fluctuations in two smallest CH(Weeks and Thurston) models are investigated. As shown in Sec.~II, it is numerically shown that the expansion coefficients of the two CH spaces behave as if they are random Gaussian numbers at almost all the places. If one recognizes the Laplacian as the Hamiltonian of a free particle, each eigenmode is interpreted as a wavefunction in a stationary state. The observed behavior is consistent with a prediction of RMT which has been considered to be a good empirical theory that describe the statistical properties of quantum mechanical systems whose classical counterparts are strongly chaotic. However, as we have seen, the global symmetries in the system can veil the generic properties. For instance, some eigenmodes on the Thurston manifold have a $Z2$ symmetry at a point where the injectivity radius is maximal. For these eigenmodes, the expansion coefficients are strongly correlated;hence they can no longer considered to be random Gaussian numbers. \\\\ \\indent Because the eigenmodes actually satisfy the periodic boundary conditions, there are points on a sphere $S^2$ which are identified with different points on $S^2$. These points form pairs of circles which are identified by the periodic boundary conditions \\cite{Cornish4}. If one could identify all the circles on a sphere, one would be able to construct the corresponding CH space \\cite{Weeks}. Similarly, if one could identify all the fixed points and the corresponding symmetries, one would be able to construct a CH space which have these symmetries. The observed ``randomness'' in the eigenmodes is actually determined by these simple structures. \\\\ \\indent In order to understand the symmetric structures of the CH spaces, it is useful to choose an observing point (base point) at which one enjoys symmetries as many as possible. However, in reality, there is no natural reason to consider fluctuations at only these particular points since the CH spaces are globally inhomogeneous. \\\\ \\indent Since the CMB fluctuations can be written in terms of a linear combination of eigenmodes, the fluctuations in CH models are almost spatially ``isotropic'' if averaged all over the space except for very limited places at which the eigenmodes have certain symmetries provided that the eigenmodes are Gaussian. The spatial ``isotropy'' implies that the contribution of non-diagonal terms in the two-point correlation functions are negligible.  Thus the validity of the statistical tests using the angular power spectrum $C_l$ \\cite{Aurich3,Inoue2,Inoue3,Cornish1} cannot be questioned on the ground that the background space is anisotropic at a certain point. \\\\ \\indent If one assumes that the initial fluctuations are Gaussian as in the standard inflationary scenarios, the temperature fluctuations are described by isotropic non-Gaussian random fields since they are written in terms of a sum of products of two independent random Gaussian variables, namely the initial perturbations and the expansion coefficients of the eigenmodes. The distribution functions of the expansion coefficients $b_{lm}$ for the sky maps at large values are slowly converged to zero than the Gaussian distribution with the same variance and the cosmic variances are found to be larger than that of the Gaussian models. \\\\ \\indent The increase in the variances are much conspicuous for topological quantities at large or small threshold levels. On the other hand, the mean behavior is well approximated by the Gaussian predictions. Therefore, the obtained results agree with the COBE DMR 4-year maps analyzed in\\cite{Colley,Kogut}. In real observations one has to tackle with what obscure the real signals such as pixel noises, galactic contaminations, beam-smoothing effect and systematic calibration errors which have not been considered in this paper. The absence of large deviations from the mean values at large or small threshold levels in the current data may be due to these effects, which will be much explored in the future work. \\\\ \\indent Although the recent observations seem to prefer the flat FRW models with the cosmological constant, the evidence is not perfectly conclusive. If one includes the cosmological constant for a fixed curvature radius, the radius of the last scattering surface (horizon) at present in unit of curvature radius becomes large. Therefore the observable imprints of the non-trivial topology of the background space become much prominent. For instance, the number $N_f$ of copies of the fundamental domains inside the last scattering at the present slice is approximately 27.9 for a Weeks model with $\\Omega_\\Lambda\\!=\\!0.6$ and $\\Omega_m\\!=0.2\\!$ whereas $N_f\\!=\\!4.3$ if $\\Omega_\\Lambda\\!=\\!0$ and $\\Omega_m\\!=0.8\\!$. Thus we have still great possibilities in detecting the non-trivial topology by the future satellite missions such as MAP and PLANCK which will provide us much better information on the statistical properties of the real signals.  The large deviations of the topological quantities from the mean values would be the good signals that indicate the hyperbolicity (negative curvature) and the finiteness (smallness) of the universe in addition to the direct observation of the periodic structures peculiar to each non-trivial topology (see \\cite{Luminet} for recent developments). \\\\ \\indent \\vspace{3cm}"
    },
    "0002/astro-ph0002138_arXiv.txt": {
        "abstract": "We present the  discovery and spectroscopic identification  of two very high proper motion ancient white dwarf stars, found in a systematic proper motion survey. Their kinematics and  apparent magnitude clearly indicate that  they are halo members,  while their optical spectra are  almost identical to  the recently identified cool  Halo  white dwarf WD0346+246.  Canonical   stellar halo models predict a white dwarf volume  density of two orders of magnitude less than  the $\\rho \\sim  7\\times 10^{-4} \\msun\\pc^{-3}$ inferred from this survey.  With the caveat that the sample size is very small, it appears that a significant fraction, $\\sim 10$\\%, of the local dark matter  halo is in the form of very old, cool, white dwarfs. ",
        "introduction": "Unless our present understanding of gravitation is incorrect, the Milky Way and most other galaxies, possess a halo of dark,  or barely luminous, matter that extends well beyond  the  visible boundaries  of the  galaxy.   Present cosmological constraints  strongly  favour non-baryonic dark  matter, though the problem  is  complicated by the  possibility that  there  may be several types of dark matter; for instance, galactic dark  matter may have little to do with dark  matter in galaxy clusters or  on  larger scales \\citet{bosma}. In our own Galaxy, this Halo component appears to extend out  to at least 50 kpc as traced by the Magellanic Clouds \\citep{lin}, and may well continue to beyond  the orbit  of  Leo~I at  $\\sim  200\\kpc$ distance  \\citep{zaritsky}. Given the large  uncertainty in the extent of  these structures, their total mass is  poorly constrained, but  even within 50  kpc it is clear  that they must contain orders of magnitude  more mass than the clearly visible baryons in the form of stars and gas.  One  of  the  few  direct  constraints  on  the  mass  distribution  of  the constituent  particles of  galactic halos  comes from  the  Magellanic Cloud microlensing experiments  \\citep{alcock97, afonso}.  The  most recent report of  the detection rate  of microlensing  events \\citep{alcock00}  supports a simple model  where approximately 20\\% of the  Milky Way halo is  made up of objects of  $\\sim 0.5\\msun$, consistent  with the result  of \\citet{afonso}. Though  the halo  fraction and  object mass  are model  dependent,  the main uncertainty  arises  from  the  possibility  that  the  LMC  may  possess  a substantial (dynamically hot) stellar halo of its own.  Given the absence of evidence for  this latter possibility, the microlensing  conclusions must be taken  seriously.  These  MACHOs  are  most  likely  situated  approximately half-way to  the LMC,  where the optical  depth to microlensing  is highest. However, unless  the MACHOs were effectively  decoupled from nucleosynthesis (as would presumably be the case for primordial black holes), it is unlikely that   MACHOs  fill   galactic   halos  entirely,   as   this  would   break nucleosynthesis constraints.  Apart from ancient  white dwarfs (WDs), all stellar  MACHO candidates can be safely  ruled out.   Ancient WDs,  which have  stellar masses  in  the range measured by the MACHO experiment ($0.1$--$1.0 \\msun$), are not yet ruled out by direct starcount  observations, but are difficult to  envisage as a MACHO population due to several indirect  constraints (see e.g., Fields, Freese \\& Graff  1998;  Graff  \\etal\\   1999),  including  most  importantly  chemical enrichment  of  the  early  Galaxy  \\citep{gibson}, though  for  a  possible work-around of this problem, see \\citet{chabrier}.  This renewed interest in ancient  Halo WDs  is mainly  due  to a  theoretical reappraisal  of the  WD cooling function.  New models  with a self-consistent treatment of radiation transfer    and    the    inclusion    of   molecular    hydrogen    opacity \\citep{hansen,saumon}   drastically  change   the   predicted  colours   and magnitudes of the  oldest, and hence coolest, WDs.   In particular, hydrogen atmosphere (DA)  WDs with  ages $\\simgt 10\\Gyr$  have depressed red  and NIR flux; contrary to previous expectations  they become bluer in, for instance, ${\\rm V-I}$, with age.  \\citet{ibata} recently  presented a proper motion analysis  using two epochs of I-band  imaging data in the  Hubble Deep Field.  They  identified a small sample 2--5 faint, and relatively  blue compact sources, that appear to move with respect to the numerous background galaxies.  The colours, motions, and number of  these sources  are consistent with  the hypothesis that  they are members of a numerous halo population of ancient WDs, with a spectral energy distribution similar to that predicted by the Hansen models.  The local mass density of this population appears  similar to that of the ``standard'' halo model   used  by   the  MACHO   and  EROS   collaborations  \\citet{alcock97, palanque-delabrouille}.  Unfortunately,  the measured proper  motions are at the detection limit, and they are  exceedingly faint (${\\rm V \\sim 29}$), so spectroscopic confirmation is currently impossible.  Striking confirmation that  the new WD models  are more appropriate has been given  by \\citet{hodgkin}.    Their  spectrum of WD0346+246   shows  a large supression of the NIR flux relative to a black-body model fit to the visible radiation, as predicted by \\citet{hansen}  for ancient DA WDs.  Furthermore, the proper motion of WD0346+246 and the  measured parallax, indicate that it is most likely a Halo member.  Therefore,  Halo  WDs do  seem to  exist and  the DA  subset  have bluer colours  than originally  expected (WD0346+246  has ${\\rm  V-I=1.52}$).  The next  important  question to  address  is the  local  mass  density of  this population,  in  particular,  are  they  present in  sufficient  numbers  to contribute sgnificantly to the proposed Halo of MACHOs? ",
        "conclusions": "No ancient  WDs were found  in survey  3.  This is  partly due to  the added incompleteness  from  the  requirement  that  a star  be  detected,  and  be classified as stellar,  on 3, rather than 2, plates.  At  each epoch, for $R \\sim 19$, the  probability of detection and correct  classification is $\\sim 85$\\%. For  rapidly moving sources, there  is also a $\\sim  5$\\% chance that they  will randomly  be situated  close (within  $3\\scnp$) to  the detection isophote of a  stationary source, and so not be detected  as a point source. We  estimate the  completeness of  surveys~1 and  2 is,  respectively, $\\sim 65$\\% and $\\sim 80$\\%, while  survey~3 is $\\sim 50$\\% complete.  Survey~3 is also  not sensitive to  faint ${\\rm  R>18}$, high  PM $\\mu>1.5\\scnp/\\yr$ stars: one of  the two detected HPM stars is outside  of the selection limit of this survey.  This implies that  the efficiency of the detection is close to  half  the   efficiency  of  the  first  two   strategies,  reducing  the completeness of survey~3 to $\\sim 25$\\%.  Discounting the overlapping regions between   adjacent plates, and the  area occupied by bright stars  and unusable regions of the  plates, surveys  1, 2 and 3 have, respectively, areas of $551\\Box^{\\circ}$, $377\\Box^{\\circ}$, and $506\\Box^{\\circ}$,    giving a  total   effective    area  of  $\\Omega  \\sim 790\\Box^{\\circ}$ explored.  We  have only  followed up  obvious detections  up to  ${\\rm  R=19}$.  Model predictions (B. Hansen, priv.  comm.)   for $T>3500$~K DA WDs are that ${\\rm M_R<16.4}$, so we are sensitive  to a distance modulus of ${\\rm m-M_0=2.7}$, that is, a distance of $d=33\\pc$.  The finding of $n=2$ stars closer than $d=33\\pc$,  over the $\\Omega \\sim 790 \\Box^\\circ$ area of the survey implies a local density of: \\begin{equation} \\rho \\sim n {{3} \\over {\\Omega d^{3}}}  M_\\star \\, \\end{equation} where $M_\\star$ are the individual object masses. Lacking any information on individual  masses, we assume  a reasonable  value  of $M_\\star \\sim 1\\msun$ \\citep{chabrier}.  Thus the local mass density of this population is: \\begin{equation} \\rho \\sim 0.0007 \\msun\\pc^{-3} \\, . \\end{equation} The expected mass density of white dwarfs in the stellar spheroid, given the subdwarf  star   counts,  and  assuming   a  standard  IMF,   is  $1.3\\times 10^{-5}\\msun \\pc^{-3}$  \\citep{gould}, that is, approximately  two orders of magnitude  lower than  this estimate.   This suggests  that the  IMF  is not universal,  and was substantially  steeper in  the progenitor  population of these WDs.  The  density derived  above is  $\\sim 10$\\%  of the  local  density ($0.0079 \\msun\\pc^{-3}$) of the ``standard'' dark matter halo model used by the MACHO collaboration \\citep{alcock97}.  However, the  halo mass fraction in WDs could be even larger, as  surveys 2 and 3 are probably not  sensitive to the older (and bluer) DA  WDs; while similar aged helium atmosphere  (DB) WDs would be too faint to be detected by any  of these methods (see Hansen 1998). We note that our  density estimate is consistent  with the updated  results from the MACHO collaboration \\citep{alcock00}, which appeared after the submission of this paper.  Thus it appears  that a significant fraction of galactic  dark matter may be baryonic, and in the form of cool and hence very old WDs, the remants of the first stellar populations.   We caution that this analysis  is based on very low  number statistics,  and  therefore requires  confirmation.  An  all-sky survey to ${\\rm R \\sim 19}$ at Galactic latitudes $b > |30|$ (as is feasable to undertake with extant photographic plates), and with completeness similar to the  present work, should find  $\\sim 14$ such stars,  providing a sample with which the Halo may be  age-dated in a way independent of isochrone fits \\citep{richer}.  Also, very deep HST proper motion data towards the Galactic center and anticenter directions  will constrain the radial density gradient of  this population,  thereby  checking  whether this  is  a local,  perhaps transient, feature, or a massive Galactic structure.  \\begin{figure} \\includegraphics[width=\\hsize]{WD_Halo.fig01.ps} \\figcaption{The  flux-calibrated  spectra  of  the  two  VHPM  stars  (left: F351-50,  right: F821-07).  The  black body  SED models,  fit to  $\\lambda < 6500\\Aa$) have  been overlaid.  The  black body temperatures are  3500~K and 4100~K for, respectively, F351-50 and F821-07.} \\end{figure}  \\begin{figure} \\includegraphics[width=\\hsize]{WD_Halo.fig02.ps}  \\figcaption{Plate  direct  images   ($2\\mcnd5\\times2\\mcnd5$)  of  (top  row) F351-50 (from  left to right: ESO B,  UKST ${\\rm B_J}$, UKST  R) and (bottom row)  F821-07 (from  left to  right: Palomar  O, UKST  ${\\rm B_J}$,  UKST R) showing  proper motion.} \\end{figure}"
    },
    "0002/astro-ph0002412_arXiv.txt": {
        "abstract": "If Gamma-Ray-Bursts (GRBs) occur at high redshifts, then their bright afterglow emission can be used to probe the ionization and metal enrichment histories of the intervening intergalactic medium during the epoch of reionization.  In contrast to other sources, such as galaxies or quasars, which fade rapidly with increasing redshift, the observed infrared flux from a GRB afterglow at a fixed observed age is only a weak function of its redshift.  This results from a combination of the spectral slope of GRB afterglows and the time-stretching of their evolution in the observer's frame.  Assuming that the GRB rate is proportional to the star formation rate and that the characteristic energy output of GRBs is $\\sim 10^{52}~{\\rm ergs}$, we predict that there are always $\\sim 15$ GRBs from redshifts $z\\ga 5$ across the sky which are brighter than $\\sim 100$ nJy at an observed wavelength of $\\sim 2\\mu$m.  The infrared spectrum of these sources could be taken with the future {\\it Next Generation Space Telescope}, as a follow-up on their early X-ray localization with the {\\it Swift} satellite. ",
        "introduction": "The past decade has been marked by major observational breakthroughs concerning the properties of the Gamma Ray Burst (GRB) sources.  The Burst and Transient Source Experiment (BATSE) on board the Compton Gamma Ray Observatory (Meegan \\etal 1992) showed that the GRB population is distributed isotropically across the sky, and that there is a deficiency of faint GRBs relative to a Euclidean distribution.  These were the first observational clues indicating a cosmological distance scale for GRBs.  The localization of GRBs by X-ray observations with the BeppoSAX satellite (Costa \\etal 1997) allowed the detection of afterglow emission at optical (e.g., van Paradijs \\etal 1997) and radio (e.g., Frail \\etal 1997; Frail, Waxman \\& Kulkarni 1999) wavelengths up to more than a year following the events (Fruchter et al. 1999; Frail, Waxman \\& Kulkarni 1999).  The afterglow emission is characterized by a broken power-law spectrum with a peak frequency that declines with time.  The radiation can be well modeled as synchrotron emission from a decelerating blast wave, created by the GRB explosion in an ambient medium, plausibly the interstellar medium of galaxies (Waxman 1997; Wijers \\& Galama 1999; M\\'esz\\'aros 1999; but see also Chevalier \\& Li 1999).  The detection of spectral features, such as metal absorption lines (Metzger \\etal 1997), in some optical afterglows allowed an unambiguous identification of the cosmological distance scale. By now, the redshift of almost a dozen GRBs has been identified either through the detection of absorption features in the afterglow spectra or of emission lines from host galaxies (see Fig. 3 in Kulkarni et al. 2000).  The central engine of GRBs is still unknown. Since the inferred energy release is comparable to or higher than that in supernovae, most popular models relate GRBs to stellar remnants, such as neutron stars or black holes (see, e.g., Eichler \\etal 1989; Paczy\\'{n}ski 1991; Usov 1992; Mochkovitch \\etal 1993; Paczy\\'{n}ski 1998; MacFadyen \\& Woosley 1999). Recently it has been claimed that the late evolution of some rapidly declining optical afterglows shows a component which is possibly associated with supernova emission (e.g., Bloom \\etal 1999; Reichart 1999). If the supernova association will be confirmed by detailed spectra of future afterglows, the GRB phenomenon will be linked to the terminal evolution of massive stars.  Any association of GRBs with the formation of single compact stars implies that the GRB rate should trace the star formation history of the universe (Totani 1997; Sahu \\etal 1997; Wijers \\etal 1998; but see Krumholz, Thorsett \\& Harrison 1998).  Owing to their high brightness, GRB afterglows might be detected out to exceedingly high redshifts. Similarly to quasars, the broad band emission of GRB afterglows can be used to probe the absorption properties of the intergalactic medium (IGM) out to the epoch when it was reionized at a redshift $z\\sim 10$ (Miralda-Escud\\'e \\& Rees 1998; Loeb 1999).  Lamb \\& Reichart (1999) have extrapolated the observed gamma-ray and afterglow spectra of known GRBs to high redshifts and emphasized the important role that their detection might play in probing the IGM.  Simple scaling of the long-wavelength spectra and temporal evolution of afterglows with redshift implies that at a fixed time lag after the GRB in the observer's frame, there is only mild change in the {\\it observed} flux at infrared or radio wavelengths with increasing redshift. This results in part from the fact that afterglows are brighter at earlier times, and that a given observed time refers to an earlier intrinsic time in the source frame as the source redshift increases. The mild dependence of the long-wavelength flux on redshift is in contrast with other high-redshift sources such as galaxies or quasars, which fade rapidly with increasing redshift (Haiman \\& Loeb 1998; 1999).  The ``apparent brightening'' of GRB afterglows with redshift could be further enhanced by the expected increase in the mean density of the interstellar medium of galaxies at increasing redshifts (Wood \\& Loeb 1999).  It therefore appears natural to use GRBs as an important tool in probing the high-redshift universe and its star formation history (Blain \\& Natarajan 1999).  Since GRBs are rare, all-sky searches for their early $\\gamma$-ray emission are needed before follow-up observations at much longer wavelength are conducted.  In this paper we model the emission from GRB afterglows and follow their number counts at different wavelengths as a function of redshift. For simplicity, we assume that the ambient gas surrounding GRB sources is the interstellar medium of their host galaxies.  In \\S 2 we review our model for both the relativistic and non-relativistic stages of the afterglow emission, as well as a simple prescription for the redshift evolution of the interstellar medium of galaxies.  The numerical results and their dependence on model assumptions are discussed in \\S \\ref{res}. Finally, \\S \\ref{disc} summarizes our main conclusions and describes the implications of our results. ",
        "conclusions": "\\label{disc}  We have calculated the expected fluxes and number counts of high-redshift GRB afterglows, assuming that their rate of occurrence is proportional to the star formation rate. We have computed the observed properties of GRB lightcurves at different wavelengths ranging from the optical to the radio, treating both the relativistic and sub-relativistic phases of the expansion of their remnants.  Our main result is that at a fixed observed time after the GRB event, the characteristic afterglow flux is not decreasing rapidly with increasing redshift (Figs. 1a, 1b), in contrast with other high-redshift sources, such as galaxies or quasars.  Hence, the broad-band spectrum of GRB afterglows is ideally suited for probing the ionization state and metal content of the intergalactic medium at high redshifts, particularly during the epoch of reionization (Loeb 1999; Lamb \\& Reichart 1999).  The only difficulty in using GRBs as probes of the high-redshift universe is that they are rare, and hence their detection requires surveys which cover a wide area of the sky (see the vertical axis in Figures~\\ref{fig2}--\\ref{fig4}).  The simplest strategy for identifying high-redshift afterglows is through all-sky surveys in the $\\gamma$-ray or X-ray regimes.  In particular, detection of high redshift sources will become feasible with the high trigger rate provided by the forthcoming {\\it Swift} satellite, to be launched in 2003.  {\\it Swift} is expected to observe and localize $\\sim$300 GRBs per year, and to repoint within 20-70 seconds its on-board x-ray and UV-optical instrumentation for continued afterglow studies. The high-resolution GRB coordinates obtained by {\\it Swift}, will be transmitted to Earth within $\\sim$50 seconds. Deep follow-up observations will then become feasible from the ground or using the highly-sensitive infrared instruments on board the {\\it Next Generation Space Telescope} (NGST), scheduled for launch in 2008.  {\\it Swift} is sufficiently sensitive to trigger on the $\\gamma$-ray emission from GRBs at redshifts $z\\ga 10$ (Lamb \\& Reichart 1999).  We note that even if some GRBs occur outside of galaxies, the intergalactic medium would be sufficiently dense to produce an afterglow at these high-redshifts since its mean density is $\\ga 10^{-4} [(1+z)/10]^3~{\\rm cm^{-3}}$.  The effect of beaming is expected to be minor for sources with millimeter and radio fluxes $\\la 10^{-5}$ Jy, since the number counts at these fluxes are dominated by sources which are in the sub-relativistic phase of their expansion, and for which the energy output in the explosion is already isotropized (see Fig.~\\ref{fig6}b,~\\ref{fig6}c).  For a characteristic energy output of GRBs of $\\sim 10^{52}$ ergs, our model implies that at any time there should be $\\sim 15$ GRBs from redshifts $z \\simgt 5$ across the sky, which are brighter than $100$ nJy at an observed wavelength of 2 $\\mu$m.  The {\\it Next Generation Space Telescope} (NGST) will be able to measure the spectra of these sources. Prior to reionization, the spectrum of GRB afterglows might reveal the long sought-after Gunn-Peterson trough (Gunn \\& Peterson 1965), due to absorption by the neutral intergalactic medium."
    },
    "0002/astro-ph0002355_arXiv.txt": {
        "abstract": "The radiation we observe from blazars is most likely the product of the transformation of bulk kinetic energy into random energy. This process must have a relatively small efficiency (e.g. 10\\%) if jets are to power the extended radio--structures. Recent results suggest that the average power reaching the extended radio regions and lobes is of the same order of that produced by accretion and illuminating the emission line clouds. Most of the radiative power is produced in a well localized region of the jet, and, at least during flares, is mainly emitted in the $\\gamma$--ray band. A possible scenario qualitatively accounting for these facts is the internal shock model, in which the central engine produces a relativistic plasma flow in an intermittent way. ",
        "introduction": "We believe that the continuum radiation we see from blazars comes from the transformation of bulk kinetic energy, and possibly Poynting flux, into random energy of particles, which quickly produce beamed emission through the synchrotron and the inverse Compton process. This is analogous to what we believe is happening in gamma--ray bursts, although the bulk Lorentz factor of their flow is initially larger.  Evidences for bulk motion in blazars with Lorentz factors between 5 and 20 have been accumulated along the years, especially through the monitoring of superluminally moving blobs on the VLBI scale (Vermeulen \\& Cohen 1994), and, more recently, through the detection of very large variable powers emitted above 100 MeV (see the third EGRET catalogue, Hartman et al., 1999), which require beaming for the source to be transparent to photon--photon absorption (e.g. Dondi \\& Ghisellini, 1995).  The explanation of intraday variations of the radio flux, leading to brightness temperatures in excess of $T_{\\rm B}=10^{18}$ K (much exceeding the Compton limit) are instead still controversial (Wagner \\& Witzel 1995). Interstellar scintillation is surely involved, but it can work only if the angular diameter of the variable sources is so small to nevertheless lead to $T_{\\rm B}=10^{15}$ K, which requires either a coherent process to be at work (e.g. Benford \\& Lesch 1998) or a Doppler factor of the order of a thousand.  Another controversial issue is the matter content of jets. We still do not know if they are dominated by electron--positron pairs or by normal electron--proton plasma (see the reviews by Celotti, 1997, 1998).  Part of our ignorance comes from the difficulty of estimating intrinsic quantities, such as the magnetic field and the particle densities, using the observed flux, which is strongly modified by the effects of relativistic aberration, time contraction and blueshift, all dependent on the unknown plasma bulk velocity and viewing angle. Furthermore it is now clear (especially thanks to multiwavelength campaigns) that the blazar phenomenon is complex.  On the optimistic side, we have for the first time a complete information of the blazar energy output, after the discovery of their $\\gamma$--ray emission, and some hints on the acceleration process, through the behaviour of flux variability detected simultaneously in different bands (see the review by Ulrich, Maraschi \\& Urry 1996). Also, blazar research can now take advantage of the explosion of studies regarding gamma--ray bursts, which face the same problem of how to transform ordered to random energy to produce beamed radiation (for reviews: Piran 1999; Meszaros 1999). ",
        "conclusions": "Here I will dare to assemble different pieces of information gathered in recent years in a coherent, albeit still preliminary, picture.  There is a link between the extraction of gravitational energy in an accretion disk and the formation and acceleration of jets, since both have the same power. Objects of low luminosity accretion disks also lack strong emission lines, suggesting that it is the paucity of ionizing photons, not of gas, the reason for the lack of strong lines in BL Lacs. Correspondingly, this implies that, if FR I are the parents of BL Lacs, they also have intrinsically weak line emission (i.e. no need for an obscuring torus). Despite the fact that the jet power in blazars spans at least four orders of magnitude, the average bulk Lorentz factor is almost the same, suggesting a link between the power and the mass outflowing rate: their ratio is constant. In the region where most of the radiation is produced, the jet is heavy, in the sense that protons carry most of the bulk kinetic energy. There the jet dissipates $\\sim 10\\%$ of its power and produces beamed radiation. The power dissipated at larger distances is much less, and therefore the jet can transport $\\sim 90\\%$ of its original power to the radio extended regions. One way to achieve this is through internal shocks, which can explain why the major dissipation occurs at a few hundreds Schwarzchild radii, why the efficiency is of the order of 10\\%, and give clues on the observed variability timescales and even on why electrons are accelerated at a preferred energy. The spectral energy distribution of blazars depends on where shell--shell collisions take place, and on the amount of seed photons present there. Even in a single source it is possible that the separation of two consecutive shells is sometimes large, resulting in a collision occuring outside the broad line region. In this case the corresponding spectrum should be produced by the synchrotron self--Compton process only, without the contribution of external photons: we then expect a simultaneous optical--$\\gamma$--ray flare of roughly equal powers (but with the self--Compton flux varying quadratically, see Ghisellini \\& Maraschi 1966). This is what should always happen in lineless BL Lac objects. On the other hand, if the initial separation of the two shells is small (or the $\\Gamma$--factor of the slower one is small), the collision takes place close to the disk. X--rays produced by the disk would then absorb all the produced $\\gamma$--rays and a pair cascade would develop, reprocessing the power originally in the $\\gamma$--ray band mainly into the X--ray band. We should therefore see an X--ray flare without accompanying emission above $\\Gamma m_{\\rm e} c^2$.   Pairs of shells which have already collided can interact again between themselves, at distances appropriate for the radio emission. This offers the interesting possibility to explain why the radio luminosity is related with the $\\gamma$--ray one, and why radio flares are associated with flares at higher frequencies. Work is in progress in order to quantitatively test this idea against observations."
    },
    "0002/astro-ph0002449_arXiv.txt": {
        "abstract": "We investigate the properties of damped Lyman-$\\alpha$ systems (DLAS) in semi-analytic models, focusing on whether the models can reproduce the kinematic properties of low-ionization metal lines described by \\nocite{pw:97,pw:98}{Prochaska} \\& {Wolfe} (1997b, 1998). We explore a variety of approaches for modelling the radial distribution of the cold neutral gas associated with the galaxies in our models, and find that our results are very sensitive to this ingredient. If we use an approach based on \\nocite{fe:80}{Fall} \\& {Efstathiou} (1980), in which the sizes of the discs are determined by conservation of angular momentum, we find that the majority of the DLAS correspond to a single galactic disc. These models generically fail to reproduce the observed distribution of velocity widths. In alternative models in which the gas discs are considerably more extended, a significant fraction of DLAS arise from lines of sight intersecting multiple gas discs in a common halo. These models produce kinematics that fit the observational data, and also seem to agree well with the results of recent hydrodynamical simulations. Thus we conclude that Cold Dark Matter based models of galaxy formation can be reconciled with the kinematic data, but only at the expense of the standard assumption that DLAS are produced by rotationally supported gas discs whose sizes are determined by conservation of angular momentum. We suggest that the distribution of cold gas at high redshift may be dominated by another process, such as tidal streaming due to mergers. ",
        "introduction": "This paper is the first in a series of papers that examines the properties of Damped Lyman-$\\alpha$ Systems (DLAS) in the context of Cold Dark Matter (CDM) based Semi-Analytic Models (SAMs). Traditionally, DLAS are believed to be the progenitors of present day spiral galaxies \\nocite{wolfe:95}({Wolfe} 1995) and thus any model of galaxy formation must also account for their properties. The current wealth of observational data on DLAS includes their number density, column density distribution, metallicities, and kinematic properties --- see \\nocite{lwt:95,stor:96,sw:00,pett:94,lu:96,pett:97, pw:97,pw:98,pw:99,pw:00,wp:00}{Lanzetta}, {Wolfe} \\& {Turnshek} (1995); {Storrie-Lombardi}, {Irwin}, \\&  {MCMahon} (1996); {Storrie-Lombardi} \\& {Wolfe} (2000); {Pettini} {et~al.} (1994); {Lu} {et~al.} (1996); {Pettini} {et~al.} (1997); {Prochaska} \\& {Wolfe} (1997b, 1998, 1999, 2000); {Wolfe} \\& {Prochaska} (2000). These data potentially provide important constraints on cosmology and theories of galaxy formation. Here we especially focus on the new kinematic data.  Previously, the number density of DLAS has been used to provide constraints on cosmological models \\nocite{mm:94,kc:94,mb:94,klyp:95}({Mo} \\& {Miralda-Escude} 1994; {Kauffmann} \\& {Charlot} 1994; {Ma} \\& {Bertschinger} 1994; {Klypin} {et~al.} 1995). These studies assumed a simple correspondance between collapsed dark matter halos and cold gas to obtain upper limits on the amount of cold gas that could be present. Gas cooling, star formation, supernovae feedback, and ionization were neglected. A different approach was used by \\nocite{lwt:95,wlfc:95,pf:95,pfh:99}{Lanzetta} {et~al.} (1995); {Wolfe} {et~al.} (1995); {Pei} \\& {Fall} (1995); {Pei}, {Fall} \\& {Hauser} (1999), in which the observed metallicities and {\\emph{observed}} number densities of the DLAS were used to model global star formation and chemical enrichment in a self-consistent way. The latter approach was set in a classical ``closed-box'' style framework rather than a cosmological context.  Clearly, in order to model DLAS realistically one needs to include the astrophysical processes of gas dynamics and cooling, star formation, and chemical enrichment within a cosmological framework. However, this is a challenge with our current theoretical and numerical capability. Cosmological $N$-body simulations with hydrodynamics are hampered by the usual limitations of volume and resolution. This is apparent in, for example, the recent work by \\nocite{gard:99}{Gardner} {et~al.} (1999), in which it was found that even rather high-resolution hydrodynamical simulations could not account for most of the observed DLAS. \\nocite{gard:99}{Gardner} {et~al.} (1999) concluded that the majority of the damped Ly-$\\alpha$ absorption must arise from structures below the resolution of their simulations. In addition, it is well known that such simulations fail to reproduce the sizes and angular momenta of present day observed spiral galaxies \\nocite{stei:99}({Steinmetz} 1999). One might therefore be suspicious of the accuracy of their representation of the spatial distribution of the cold gas that gives rise to DLAS at high redshift. Because observational samples of DLAS are cross-section weighted, these properties are likely to introduce crucial selection effects. Semi-analytic approaches can deal with nearly arbitrary resolution and volumes, but are limited in the sophistication and accuracy of their physical ``recipes''. In particular, most previous SAMs have focussed on the bulk properties of galaxies, and have not attempted to model the spatial location of galaxies relative to one another or the spatial distribution of gas and stars within galaxies.  The only previous attempt to model the properties of DLAS in a CDM framework is the work of \\nocite{kauf:96}{Kauffmann} (1996) (hereafter K96).  In K96 the radial distribution of cold gas in galactic discs was modelled by assuming that the initial angular momentum of the gas matched that of the halo, and that angular momentum was conserved during the collapse. Star formation was then modelled using the empirical law of \\nocite{kenn:89,kenn:98}{Kennicutt} (1989, 1998), in which the star formation is a function of the surface density of the gas, and cuts off below a critical threshold density. K96 then showed that the number density, column density distribution, and metallicities of observed DLAS could be reasonably well reproduced within the Standard Cold Dark Matter (SCDM) cosmology, and predicted the distribution of circular velocities of discs that would give rise to DLAS. Assuming that each observed DLAS corresponds to a single galactic disc, this can then be compared with the observed distribution of velocity widths derived from the kinematics of unsaturated, low-ionization metal lines \\nocite{pw:97,pw:98}({Prochaska} \\& {Wolfe} 1997b, 1998).  \\nocite{pw:97}{Prochaska} \\& {Wolfe} found the velocity distribution predicted by K96 to be strongly inconsistent with their data. Furthermore \\nocite{jp:98}{Jedamzik} \\& {Prochaska} (1998) showed that the thick rotating disc model favored by \\nocite{pw:97}{Prochaska} \\& {Wolfe} (1997b) could only be reconciled with a finely tuned CDM model. But CDM actually predicts that halos will have much substructure, and \\nocite{hsr:98}{Haehnelt}, {Steinmetz} \\&  {Rauch} (1998) found that large $\\delv$ velocity profiles consistent with those observed by \\nocite{pw:97}{Prochaska} \\& {Wolfe} are produced in their very high-resolution hydrodynamical simulations. These profiles arose not from the rotation of a single disc, but from lines of sight intersecting multiple proto-galactic ``clumps''. Subsequently, \\nocite{mm:99}{McDonald} \\& {Miralda-Escud\\'{e}} (1999) also showed with a simple analytical model that DLAS produced by intersection with a few gas clouds could create kinematics consistent with the observations in a CDM universe.  These results were encouraging but remain somewhat inconclusive. The hydro simulations do not allow the construction of a statistical, cross-section selected sample of DLAS, so it is difficult to assess how typical are the systems that they identified. In addition, these simulations were restricted to a single cosmology (SCDM), and did not include star formation or supernovae feedback. The generic difficulty of hydro simulations in producing reasonable discs at low redshift has already been noted. Therefore a further investigation using detailed semi-analytic models is worthwhile.  In the standard CDM picture of galaxy formation \\nocite{wr:78,bfpr:84}(based on {White} \\& {Rees} 1978; {Blumenthal} {et~al.} 1984) gas is heated to the virial temperature when a halo forms and then cools and falls into the centre of the halo where it subsequently forms stars.  In SAMs, which include the hierarchical formation of structure, this process happens numerous times as halos continually merge and form larger structures.  This naturally results in halos that may contain many gaseous discs, each one associated with a sub-halo that prior to merging had been an independent halo. In this paper, we explore the possibility that such a scenario can account for the observed kinematics of the DLAS in a manner analogous to the proto-galactic clumps of \\nocite{hsr:98}{Haehnelt} {et~al.} (1998) and the gas clouds of \\nocite{mm:99}{McDonald} \\& {Miralda-Escud\\'{e}} (1999). Here, however, the number densities, gas contents, and metallicities of these proto-galaxies are determined by the full machinery of the SAMs, which have been tuned to produce good agreement with the optical properties of galaxies at low and high redshift. We introduce new ingredients to describe the kinematics of satellite galaxies within dark matter halos, and the spatial distribution of cold gas in discs. We also include a model that is not based on the machinery of the SAMs to demonstrate that our general conclusions are not overly dependent on the specifics of how these processes are handled in the SAMs.  We start with a review of the the observational properties of DLAS (section \\ref{observations}). Next, section \\ref{models} gives a brief description of the ingredients of the SAMs, and describes how we simulate the observational selection process for DLAS and produce simulated velocity profiles.  We demonstrate in section \\ref{bad_models} that gaseous discs with sizes determined by conservation of angular momentum fail to match the kinematic data, and then in section \\ref{bigdiscs} show that acceptable solutions can be found if the gaseous discs have a large radial extent.  Section \\ref{depend} examines the sensitivity of our results to a number of model parameters.  In section \\ref{compare} we discuss the properties of the gas discs in our model and compare them to HI observations of local spirals and with the results of hydro simulations. Lastly we close with some discussion and conclusions. ",
        "conclusions": "\\label{conclusions}  We have explored the properties of DLAS in semi-analytic models of galaxy formation. These models produce good agreement with many optical properties of galaxies at low and high redshift, and the total mass of cold gas at redshift $\\sim 3$ is also in reasonable agreement with observations. It is therefore interesting to ask whether the kinematic properties, metallicities, and column densities of DLAS in these models are in agreement with observations. We investigated the dependences of these properties on cosmology, the distribution of satellite orbits, and gaseous disc scale height, and found that our results were not sensitive to these assumptions. Our results are \\emph{extremely sensitive} to our assumptions about the radial distribution of cold gas within galactic discs. Given that one believes the other components of our model, one can then perhaps learn about the distribution of cold neutral gas at high redshift. \\begin{figure} % \\centerline{\\epsfig{file=fig13.ps,width=\\linewidth}} \\vskip .5 cm \\caption{The relationship between the velocity width observed and the circular velocity of the halo in our toy model. }\\label{fdvvc} \\end{figure}  Currently popular theories of disc formation posit that the radial size of a galactic disc is determined by the initial specific angular momentum of the dark matter halo in which it forms, and that the cold gas traces the stellar component. Often, the profile of the disc is assumed to have an exponential form. We investigate several variants of such models, based on ideas in the literature such as \\nocite{fe:80}{Fall} \\& {Efstathiou} (1980), \\nocite{mmw:98}{Mo} {et~al.} (1998) and \\nocite{kauf:96}{Kauffmann} (1996). We find that the kinematics of DLAS arising in such models are in strong conflict with the observations of \\nocite{pw:97,pw:98}{Prochaska} \\& {Wolfe} (1997b, 1998). This is consistent with the previous work of \\nocite{pw:97}{Prochaska} \\& {Wolfe}, in which it was shown that if the $\\delv$ of each DLAS arises from a single rotating disc, generic CDM models can be ruled out at high confidence level. Our work has shown that in theories of disc formation based on angular momentum conservation, the resulting gaseous discs are so small that most DLAS are produced by a single disc and thus the models suffer from the familiar difficulties with the observed kinematics. In addition, although the total mass of cold gas in the models is in agreement with the estimate of $\\Omega_{\\rm gas}$ from \\nocite{stor:96}{Storrie-Lombardi} {et~al.} (1996), when we use realistic cross-section-weighted column-density criteria to select DLAS in our models, we find that the overall number density of damped systems is too small. This again seems to indicate that the covering factor of the gas is too small.  We therefore abandon the standard picture of discs and investigate a toy model in which the gas is distributed according to a Mestel distribution with a fixed truncation radius. We adjust the truncation radius as a free parameter to find the best fit with the observations, and find that the best-fit value is consistent with the expected ionization edge due to a cosmic ionizing background. This results in gas discs which are considerably more radially extended than the standard models discussed above. For this class of models, we find good agreement with the four diagnostic statistics of PW97 which describe the kinematics; in particular, the distribution of velocity widths $\\delv$ in the models now has a tail to large $\\delv \\sim 200-300 \\kms$ as in the observations. In the models, the majority of these large $\\delv$ systems arise from ``multiple hits'', lines of sight that pass through more than one rotating disc, as in the picture proposed by \\nocite{hsr:98}{Haehnelt} {et~al.} (1998). The column density distribution and metallicities of the DLAS in the models are also in reasonable agreement with the observations.  This working model for DLAS has many additional implications that may be tested by observations in the near future.  One interesting issue is the relationship between DLAS and the Lyman-break galaxies \\nocite{stei:96,lowe:97,stei:98}({Steidel} {et~al.} 1996; {Lowenthal} {et~al.} 1997; {Steidel} {et~al.} 1998), about which little is currently known observationally \\nocite{djor:97,mw:98}({Djorgovski} 1997; {Moller} \\& {Warren} 1998). Previous theoretical predictions used simplifed relations to estimate luminosities \\nocite{mmw:98,hsr:99}({Mo} {et~al.} 1998; {Haehnelt} {et~al.} 1999). Because the SAMs include detailed modelling of star formation-related processes as well as full stellar population synthesis, we are in a position to make much more detailed and perhaps more reliable predictions. Our model suggests that $20\\%$ of DLAS contain at least one galaxy with an $R$ magnitude brighter then 25.5 in the same dark matter halo.  The median projected distance between the DLAS and the Lyman-break galaxy is about 30 kpc \\nocite{mpsp:00}({Maller} {et~al.} 2000a).   Another interesting comparison is with the kinematics of the high ionization state elements \\nocite{wp:00}({Wolfe} \\& {Prochaska} 2000).  In the simple picture of the SAMs, these profiles would naturally be associated with the gas that has been shock heated to the virial temperature of the halo.  The hot gas is distributed spherically in the sub-halo, unlike the cold gas, which would explain why the velocity profiles of the high ions do not trace the low ions \\nocite{pw:97j}({Prochaska} \\& {Wolfe} 1997a).  However, the velocity widths of the two profiles, which are dominated by the motions between sub-halos within the same larger halo, would be related.  The kinematics and $f(N)$ distribution for absorbers below the damped limit but with column densities above the value where the disc is truncated (i.e. Lyman-limit systems) also provide an interesting test of our model.  Our modeling suggests that the incidence of these absorbers arising from cold discs should increase with the same slope as in Fig.~\\ref{ffnbd} and then turn over abruptly around a column density of $5 \\times 10^{19} \\cm2$. Below this column density these Lyman limit systems must be composed of more diffuse ionized gas. This is found in hydro simulations \\nocite{dhkw:99}({Dav\\'{e}} {et~al.} 1999) and supported by some observational evidence \\nocite{proc:99}({Prochaska} 1999). Thus observations at these column densities can directly probe whether gas discs reach the values of $N_t$ we require to explain the DLAS kinematics and $f(N)$. Lastly it is possible to explore how the properties of the DLAS evolve with redshift. The merging rate is a strong function of redshift \\nocite{kola:99,kola:00}({Kolatt} {et~al.} 1999, 2000) so we would expect the number of ``multiple hits'' and therefore the kinematics of the DLAS to be significantly different at low redshift. All these issues will be explored in greater detail in subsequent papers.  If one is really to accept our conclusion that high redshift discs have an extended Mestel-type radial profile, clearly we must develop a theory for their origin. It is possible that the standard theory of disc formation is applicable at low redshift, but that some other process dominates at higher redshift. For example, mergers are far more common at high redshift, and the gas fractions of discs are higher (cf. SPF). This may result in efficient transfer of orbital angular momentum to the gas, producing tidal tails that distribute the gas out to large radii. Locally, some interacting galaxies show extremely extended rotating HI, presumably resulting from such a mechanism \\nocite{hibb:99}({Hibbard} 1999). Another possibility is that starbursts triggered by mergers produce supernovae-driven outflows like those seen in local starbursts \\nocite{heck:99}({Heckman} 1999) that could also result in extended gas distributions. We find that a toy model in which the gas clouds have a bulk velocity equal to the rotation velocity of the disc, but in a random direction (i.e. not in a rotationally supported disc) still produces good agreement with the observed DLAS kinematics because the kinematics of our model are dominated by the motions of the multiple discs, not the kinematics within these discs.  It is worth noting that the surface densities of the gas discs in our model would be far below the critical value for star formation determined by observations \\nocite{kenn:89,kenn:98}{Kennicutt} (1989, 1998), implying that there may be very little star formation taking place in a `quiescent' mode. It is interesting that SPF found that a picture in which quiescent star formation at high redshift is very inefficient and most of the star formation occurs in merger-induced bursts provides the best explanation of the high redshift Lyman-break galaxies. Even in the extreme case in which quiescent star formation is completely switched off, we find that the starburst mode alone can easily produce the observed level of star formation at high redshift. Thus in the high redshift universe interactions between galaxies seem to play a rather prominent role in determining their gas properties and star formation histories."
    },
    "0002/astro-ph0002480_arXiv.txt": {
        "abstract": "The very wide double star CCDM~23239-5349 is an interesting study case of a pulsating star within a common origin pair or wide binary. {\\it Both} stars are located in the $\\delta$ Scuti instability strip. Based on data obtained at La Silla (Chile), at the Swiss 0.7m and the ESO 0.5m telescopes, we found two periodicities of about 4.7 and 5.5 cycles per day (cpd) with amplitudes of 0.014 and 0.011 mag, respectively, for HD~220392, the brightest member of the system. The most dominant periodicity was also detected in the Hipparcos Epoch Photometry data. A similar period search on the (smaller) dataset obtained for the 1 mag fainter B-component, HD~220391, however shows no periodic behaviour with an amplitude significantly above the noise level of the data (about 0.006~mag). ",
        "introduction": "\\footnotetext[1]{ Based on observations done at La Silla (ESO, Chile) and on data obtained by the Hipparcos astrometry satellite.}  The double star CCDM~23239-5349 is a wide visual system consisting of two bright stars (with $\\Delta${\\rm m} $\\simeq$ 1 mag and an angular separation of 26.5 arcsec) having similar proper motions as well as compatible parallaxes. This classifies it as a wide binary (see Sect.4.1).  Regular short-period light variations on a time scale of $\\simeq$ 5 hr have been detected for the brightest component of the system (Lampens~1992). The detailed investigation of the difference in variability and physical properties between two components of a stellar pair is particularly interesting when both companions are located in the same area of the colour-magnitude diagram: in this case both stars are situated in the $\\delta$ Scuti instability strip. The aim of this study is to search for clues to understand which factors determine the pulsation characteristics such as modes and amplitudes among $\\delta$ Scuti stars. A more extensive discussion will appear elsewhere (Lampens, Van Camp, \\& Sinachopoulos~2000). ",
        "conclusions": "What factors actually determine the pulsation characteristics, i.e., the amplitudes and modes of pulsation in the $\\delta$ Scuti instability strip? This question cannot be addressed on the basis of a single case. But it could be approached as illustrated here. Binary systems with pulsating variable components offer a unique opportunity of coupling the information obtained by astrometric means (association type - parallax - total mass) to the astrophysical quantities (luminosity ratio - colours - pulsation characteristics). For example, the detailed investigation of differences in variability between the components of a binary may provide relevant clues with respect to their pulsation characteristics. Differences in origin and age can be ruled out as well as differences in overall chemical composition. In this case the short-period pulsating component is easily identified and the information obtained on the variability can be coupled to the astrophysical parameters of {\\it each} component. It would be even better to investigate such characteristics in a visual binary for which basic information on the orbital motion can be derived.  This would allow one to obtain a direct estimation of the stellar mass, independent of any choice of modelisation. The derivation of the pulsation constant would be more straightforward too (the error on the mass defines the accuracy of Q)."
    },
    "0002/astro-ph0002213_arXiv.txt": {
        "abstract": "We investigate a model of disk galaxies  whereby viscous evolution of the gaseous disk drives material inwards to form a proto-bulge. We start from the  standard picture of disk formation through the settling of gas into a dark halo potential well, with the disk initially coming into centrifugal equilibrium with detailed conservation of angular momentum. We derive generic analytic solutions for the disk-halo system after adiabatic compression of the dark halo, with free choice of the input virialized dark halo density profile and of the specific angular momentum distribution. We derive limits on the final density profile of the halo in the central regions.   Subsequent viscous evolution of the disk is modelled by a variation of the specific angular momentum distribution of the disk, providing analytic solutions to the final disk structure.  The assumption that the viscous evolution timescale and star formation timescale are similar leads to predictions of the properties of the stellar components. Focusing on small `exponential' bulges, ones that may be formed through a disk instability, we investigate the relationship between the assumed initial conditions, such as halo `formation', or assembly, redshift $z_f$, spin parameter $\\lambda$, baryonic fraction $F$, and final disk properties such as global star formation timescale, gas fraction, and bulge-to-disk ratio. We find that the present properties of disks, such as the scale length, are compatible with a higher initial formation redshift if the re-distribution by viscous evolution is included than if it is ignored.  We also quantify the dependence of final disk properties on the ratio $F/\\lambda$, thus including the possibility that the baryonic fraction varies from galaxy to galaxy, as perhaps may be inferred from the observations. ",
        "introduction": "The current picture of disk galaxy formation and evolution has as its basis the dissipative infall of baryons within a dominant dark halo potential well (White \\& Rees 1978). The collapse and spin-up of the baryons, with angular momentum conservation, can provide an explanation for many of the observed properties of disks, with the standard initial conditions of baryonic mass fraction $F \\sim 0.1$ and dark halo angular momentum parameter $\\lambda \\sim 0.07$ (Fall \\& Efstathiou 1980; Gunn 1982; Jones \\& Wyse 1983; Dalcanton, Spergel \\& Summers 1997; Hernandez \\& Gilmore 1998; Mo, Mao \\& White 1998; van den Bosch 1998). Galaxies such as the Milky Way which have an old stellar population  in the disk, must, within the context of a hierarchical-clustering scenario, evolve through only quiescent merging/accretion, so as to avoid excessive heating and disruption of the disk (Ostriker 1990).  Further, the merging processes with significant substructure cause angular momentum transport to the outer regions, which must somehow be suppressed to allow the formation of extended disks as observed (e.g. Zurek, Quinn \\& Salmon 1988; Silk \\& Wyse 1993; Navarro \\& Steinmetz 1997).  Thus here we adopt the simplified picture that disk galaxies form from smooth gaseous collapse to centrifugal equilibrium, within a steady dark halo potential.  We discuss where appropriate below how this may be modified to take account of subsequent infall, or earlier star formation. Our model incorporates the adiabatic response of the dark halo to the disk infall, and we provide new, more general, analytic solutions for the density profile, given a wide range of initial density profiles and angular momentum distributions.  We provide new insight into the `disk-halo' conspiracy within the context of this model, demonstrating how an imperfect conspiracy is improved by the disk-halo interaction.  We explicitly include subsequent viscous evolution of the gas disk to provide the exponential profile of the stellar disk, and develop analytic expressions that illustrate the process.  The resulting radial inflow builds up the central regions of the disk and we investigate the properties of `bulges' that may form as a consequence of instabilities of the central disk.  We derive new constraints on the characteristic redshift of disk star formation. We obtain a simple relation connecting the initial conditions, such as spin parameter and baryonic mass fraction, to the efficiency of viscous evolution and star formation. ",
        "conclusions": "In the context of hierarchical clustering cosmology, the dark halo of a disk galaxy can be formed by quiescently merging small sub-halos into the primary dark halo, or by smoothly accreting matter into the dark halo. We derive the generic solution to the adiabatic infall model of disk galaxy formation pioneered by many authors (Mestel 1963; Fall \\& Efstathiou 1980; Gunn 1982; Faber 1982; Jones \\& Wyse 1983; Ryden \\& Gunn 1987; Dalcanton, Summer \\& Spergel 1997; Mo, Mao \\& White 1998; Hernandez \\& Gilmore 1998). Through exploring the allowed parameter space of dark halo profile and angular momentum distribution function, we show that the central halo density profile should be cuspy, with the power law index ranging from $-0.75 $ to $-2$, in the central regions where the disk mass dominates.   Using a modified Schmidt law of global star formation rate, we derive a simple scaling relationship between the disk gas fraction and the assembly redshift. We explicitly allow a distribution in the values of the baryonic mass fraction, $F$, in addition of the distribution in values of the spin parameter $\\lambda$. These two are found to play different role in determining the structural properties of the final disk, the star formation properties and bulge-to-disk ratio.  We mimic viscous evolution of disks by varying the specific angular momentum distribution of the disk, to redistribute angular momentum as a function of time.  We derive a consistent picture of the formation of galaxies like the Milky Way, with old stars in the disk. Under the assumption that the viscous evolution timescale is equal to the star formation timescale, we can further combine the $\\lambda$ and $F$ with the efficiency of angular momentum redistribution caused by viscosity. Assuming that small bulges are formed from their disks through bar dissolution, we can use the global bar instability condition to obtain bulge-to-total ratio, and explore the dependence on  $F$, $\\lambda$ and the viscous evolution efficiency.   The inclusion of viscous evolution has the merits of addressing several important issues:  the conspiracy between disk and halo, the formation of the exponential profile of stellar disk, the high phase space density of bulges. We have presented an analytic treatment, to illustrate these points and identify areas of particular need for more work."
    },
    "0002/astro-ph0002025_arXiv.txt": {
        "abstract": "The main goals of the KASCADE (KArlsruhe Shower Core and Array DEtector) experiment are the determination of the energy spectrum and elemental composition of the charged cosmic rays in the energy range around the knee at $\\approx 5\\,$PeV. Due to the large number of measured observables per single shower a variety of different approaches are applied to the data, preferably on an event-by-event basis. First results are presented and the influence of the high-energy interaction models underlying the analyses is discussed. \\vspace{1pc} \\vspace{-1.pc} ",
        "introduction": "The air shower experiment KASCADE \\cite{klages} aims at the investigation of the knee region of the charged cosmic rays. It is built up as a multidetector setup for measuring simultaneously a large number of observables in the different particle (electromagnetic, muonic and hadronic) components of the extended air shower (EAS). This enables to perform a multivariate multiparameter analysis for the registered EAS on an event-by-event basis to account for the non parametric, stochastic processes of the EAS development in the atmosphere. In parallel the KASCADE collaboration tries to improve the tools for the Monte Carlo simulations with the relevant physics. The code CORSIKA \\cite{cors} allows not only the detailed three dimensional simulation of the shower development in all particle components (including neutrinos) down to the observation level, but it has been implemented several high-energy interaction models. As the basic physics of these models in the relevant energy region and in the extreme forward direction cannot be tested at present days' accelerators, the test of these models emerged as one of the goals of the KASCADE experiment. The following overview is based on results presented at the 26$^{\\rm th}$ International Cosmic Ray Conference in Salt Lake City, Utah 1999 \\cite{slc}. ",
        "conclusions": "First results of the KASCADE experiment can be summarized by following statements: All secondary particle components of the showers display a kink in the size spectra. This strongly supports an astrophysical origin of the ``knee'', rather than effects of the interaction of the primaries in the atmosphere. The knee is sharper for the light primary component than for the heavy one. This result follows from the measurement as an increasing average mass of the primary cosmic rays above the observed kink, together with the energy dependent mass classification of single air showers. But none of the high-energy interaction models en vogue is able to fit the data of all observables consistently."
    },
    "0002/astro-ph0002269_arXiv.txt": {
        "abstract": "I consider analytically particle acceleration at relativistic shocks, in the limit of pitch angle diffusion and large shock Lorentz factors. The derived energy spectral index ($k = (1+\\sqrt{13})/2 \\approx 2.30$), and particle pitch angle distribution at the shock are successfully compared with the results of numerical solutions. This totally analytic derivation is completely independent of the detailed dependence of the diffusion coefficient $D(\\mu,p)$ on either parameter, and it is argued on a physical basis that the orientation of the magnetic field is also irrelevant, making $k$ a universal index. ",
        "introduction": "There is currently a growing interest in the acceleration of non--thermal particles at highly relativistic shocks. There are three classes of {\\it bona fide} relativistic sources: beyond the well--established extra--Galactic (Blazars) and Galactic (superluminal) sources, both of which exhibit superluminal motions, it is now also well-established that Gamma Ray Bursts (GRBs) display highly relativistic expansions, with Lorentz factors well in excess of $100$. Other classes of relativistic sources may include Soft Gamma Ray Repeaters (SGRs), whose recurrent explosions are largely super--Eddington, and special SNe similar to SN 1998bw, which displayed marginally Newtonian expansion ($\\approx 6\\times 10^4\\; km\\; s^{-1}$) when optical emission lines became detectable, about a month after the explosion.  With the discovery of GRBs' afterglows, it has now become feasible to derive the energy spectral index $k$ of electrons accelerated at the forward shock, as a function of the varying (decreasing) shock Lorentz factor $\\gamma$, provided simultaneous wide--band spectral coverage is available. With the launch of the USA/Italy/UK mission SWIFT, these data will become  available for a statistically significant number of bursts, testing directly models for particle acceleration at relativistic shocks. Furthermore, since GRBs must also clearly accelerate protons, the same index $k$ may determine the spectrum of ultra high energy cosmic rays observed at Earth.  However, until recently, both the lack of astrophysical motivation and the difficulty inherent in treating highly anisotropic distribution functions have stiffened research on this topic. Early work, both semi--analytic and outright numerical, has concentrated on barely relativistic flows with Lorentz factors of a few, well suited to Blazars and Galactic superluminals, but clearly insufficent for GRBs, the only exception being the numerical simulations of Bednarz and Ostrowski (1998). It is the purpose of this {\\it Letter} to perform an analytic investigation of the large $\\gamma$ limit, to establish which (if any) of the properties of the particles' distribution function depend upon the physical conditions of the fluid. ",
        "conclusions": "For ultrarelativistic particles the energy spectral index $k$ is related to $s$ by \\begin{equation} \\label{final} k = s-2 = \\frac{1 +\\sqrt{13}}{2} \\approx 2.30\\;, \\end{equation} \\begin{figure} \\epsscale{1.0} \\plotone{dist.eps} \\caption{Pitch angle distribution at the shock in the downstream frame, Eq. \\ref{incomplete} with $s$ from Eq. \\ref{ss}, with arbitrary vertical scale.} \\end{figure} which is our final result. Also, now that we know $s$, the final pitch angle distribution at the shock, but downstream, can be determined (Eq. \\ref{incomplete}), and is plotted in Fig. 2. None of these results depends upon the specific form of $D(\\mu, p)$, so that widely differing assumptions should yield precisely the same results.  How does this compare with numerical work? The near--constancy of the index $s$ (or $k$) explains moderately well the results of previous authors: Kirk and Schneider (1987) find $s= 4.3$ for their computation with the highest speed, which is however a modest Lorentz factor of $\\gamma = 5$, and a single functional dependence of $D$ on $\\mu$. Heavens and Drury (1988) find again a result of $s=4.2$ for equally modest Lorentz factors, but for two different recipes for $D$. Extensive numerical computations using a MonteCarlo technique (\\ie, totally independent of the validity of Eq. \\ref{main}) were performed by Bednarz and Ostrowski (1998), for a wide variety of assumptions about the scattering properties of the fluids. They remarked quite explicitly that the energy spectral index $k$ seemed to converge to a constant value, independent of shock Lorentz factor (provided $\\gamma \\gg 1$), magnetic field orientation angle $\\phi$, and diffusion coefficient ratio, $\\kappa_\\perp/\\kappa_\\parallel$. They found $k \\approx 2.2$. The present work confirms (for all untried forms for $D(\\mu, p)$) and extends their simulations (by yielding the exact value, and explicit forms for the particle angular distributions). The upstream angular distributions also agree well: Fig.3a of Bednarz and Ostrowski clearly shows that this is (for the highest displayed value of $\\gamma = 27$) a Dirac's delta, in agreement with the large--$\\gamma$ limit in Section 2.1. However, the downstream pitch angle distributions (in their Fig. 3b) agree well with mine, but not perfectly. Gallant, Achterberg and Kirk (1998) have claimed that there is a small error in Bednarz and Ostrowksi's distributions. As a matter of fact, my distribution (Fig. 2) agrees much better with Gallant \\etal's and Kirk and Schneider's (1987) than Bednarz and Ostrowski's, despite the very small shock Lorentz factors of these two papers ($\\gamma = 2.3$ and $\\gamma = 5$, respectively). Possibly, the small error in question may even explain the (small!) discrepancy between the two values of $k$.  A limitation applies to the claim of universality of Eqs. \\ref{up}, \\ref{incomplete} and \\ref{final}: I neglected any process altering the particles' energy during the scattering. Clearly, the results of this paper only apply in the limit $\\delta\\!p/p \\la 1$, where $\\delta\\!p$ is the typical momentum transfer in each scattering event. In the large momentum limit considered here, it seems unlikely that this constraint may be violated.  Lastly, a comment on the assumed dependence $\\propto p^{-s}$ of the distribution function upon particle momenta is in order. It can be seen from Eq. \\ref{main} that such a dependence is {\\it not} required by this equation. To see this, let us make the usual assumption that $D$ is homogeneous of degree $-r$ in $p$, \\ie, $D(\\mu, p) = q(\\mu) p^{-r}$. Then by defining a new variable $\\hat{z} \\equiv z/p^r$, we see that the form assumed by Eq. \\ref{main} after this change of variable is identical to the original one, except that now $p$ has altogether disappeared. At large $z$ (\\ie, far downstream), $f \\rightarrow f_\\infty =$ constant, and there is no $p$--dependence. This paradox is solved by noticing that the real problem to be solved involves both scattering (= Fermi acceleration) and injection. In this case, a typical injection momentum $p_0$ arises naturally, and the dimensional problem discussed above is naturally solved: we must have $f = f(...,p/p_0,...)$ where the dots indicate all other parameters. In the limit of $p_0 \\rightarrow 0$, $f$ {\\it does not} tend to a constant independent of $p_0$ as is always assumed, but tends instead to zero as $f \\rightarrow (p_0/p)^s$. Problems of this sort, though rare in astrophysics, are common in hydrodynamics, where they are called self--similar problems of the second kind (Zel'dovich 1956). They range from the deceptively simple laminar flow of an ideal fluid plast an infinite wedge (Landau and Lifshitz 1987) to the illuminating case of the filtration in an elasto--plastic porous medium (Barenblatt 1996). It is remarkable that, in the problem at hand, no such complication is necessary to fix the all--important index $s$, yet the powerful methods of intermediate asymptotics (Barenblatt 1996) and the renormalization group (Goldenfeld 1992) can be brought to bear on the intermediate $\\gamma$ cases, where no easy limiting solution can be found.  In short, what I have done in this paper is to show that the spectrum of non--thermal particles accelerated at relativistic shocks is universal, in the sense that the energy spectral index $k$, and the angular distributions in both the upstream and downstream frames (Eqs. \\ref{up}, \\ref{incomplete}, \\ref{final}, and Fig. 2) do not depend upon the scattering function $D(\\mu, p)$, the shock Lorentz factor (provided of course $\\gamma \\gg 1$), the magnetic field geometry, and the ratio of cross--field to parallel diffusion coefficients.Thus we have the result that the cosmic rays' spectra are independent of flow details in both the Newtonian (Bell 1978) and the relativistic regimes.  {}"
    },
    "0002/astro-ph0002096_arXiv.txt": {
        "abstract": "{We review constraints on models of the broad line region that are imposed by observations of the emission lines in Active Galactic Nuclei.  Comparison of  high and low ionization lines in sources with FWHM H$\\beta\\leq$4000 km s$^{-1}$ points toward low ionization line emission produced in a flattened geometry (the accretion disk?) with an associated  high ionization wind. It remains unclear whether these results can be extended to all radio quiet AGN and particularly to radio loud AGN.}  \\resumen{Revisamos las limitaciones que se pueden imponer a partir de las observaciones, sobre los modelos de la estructura de las regiones que emiten las lineas anchas en los N\\'ucleos Activos de Galaxias.  La comparaci\\'on entre  las lineas de baja- y alta ionizaci\\'on, sugiere que la emisi\\'on proviene, en parte, de una geometr\\'{\\i}a fuertemente aplanada (el disco de acreci\\'on?) y, en parte, de un viento asociado, para una fracci\\'on importante  de objetos radio callados que incluye a los nucleos Seyfert 1 de lineas angostas (NLSy1). Sin embargo, es incierto si estos resultados pueden extenderse a todos los Nucleos Activos radio callados y, sobre todo, a los radio fuertes.}   \\nonstopmode \\def\\habc{{\\sc{H}}$\\alpha_{\\rm BC}$\\/} \\def\\hbnc{{\\sc{H}}$\\beta_{\\rm NC}$\\/} \\def\\lya{{\\sc L}{\\rm y}$\\alpha$\\/} \\def\\fevii{{\\sc{Fe vii}}$\\lambda$6087\\/} \\def\\ovi{{\\sc{Ovi}}$\\lambda$1034\\/} \\def\\mgii{{\\rm Mg\\sc{ii}}$\\lambda$2800\\/} \\def\\niv{{\\sc{Niv]}}$\\lambda$1486\\/} \\def\\ciii{{\\sc{Ciii}]}$\\lambda$1909\\/} \\def\\oiiiopt{{\\sc{[Oiii]}}$\\lambda$4959,5007\\/} \\def\\o4363{{\\sc{[Oiii]}}$\\lambda$4363\\/} \\def\\oiiiuv{{\\sc{Oiii]}}$\\lambda$1663\\/} \\def\\heiiuv{{\\rm He\\sc{ii}}$\\lambda$1640} \\def\\heiiuvnc{He{\\sc{ii}}$\\lambda$1640$_{\\rm NC}$\\/} \\def\\feiiuv{\\rm Fe{\\sc ii}$_{\\rm UV}$\\/} \\def\\fei{\\rm Fe{\\sc i}} \\def\\feii{\\rm Fe{\\sc ii}} \\def\\feiii{\\rm Fe{\\sc iii}} \\def\\feiiopt{\\rm Fe{\\sc ii}$_{\\rm opt}$\\/} \\def\\feiiq{\\rm Fe{\\sc ii }$\\lambda$4570\\/} \\def\\heiiopt{\\rm He{\\sc{ii}$\\lambda$4686}\\/} \\def\\heii{He{\\sc{ii}}$\\lambda$4686\\/} \\def\\oiii{{\\sc [Oiii]}$\\lambda\\lambda$4959,5007} \\def\\oii{{\\sc [Oii]}$\\lambda$3727} \\def\\oiir{{\\sc Oi}$\\lambda$8442} \\def\\aliii{Al{\\sc  iii}$\\lambda$1860} \\def\\siiii{Si{\\sc  iii]}$\\lambda$1892} \\def\\iron{{[Fe/H]\\/}} \\def\\ironsol{{ [Fe/H]}$_{\\odot}$\\/} \\def\\feka{Fe~{\\sc  K$\\alpha$\\/}} \\def\\apj{{\\it Ap. J.}} \\def\\aap{{\\it Astr. Astroph.}} \\def\\apjs{{\\em Ap. J. Suppl.}} \\def\\nat{{\\it Nature}} \\def\\b{$\\beta$} \\def\\kms{km~s$^{-1}$} \\def\\l{$\\lambda$} \\def\\araa{{\\it Annu. Rev. Astron. Astrophys. }} \\def\\aj{{\\it Astron. J. }} \\def\\mnras{{\\it MNRAS}} \\def\\aap{{\\it Astron. Astrophys.}} \\def\\apjl{{\\it Ap. J. Lett. }} \\def\\aaps{{\\it Astron. Astrophys. Suppl. }} \\def\\REF{\\par\\noindent\\hangindent 20pt} \\def\\ab{{\\sc Z}} \\def\\absol{{\\sc Z$_{\\odot}$\\/}} \\def\\msol{$\\rm M_\\odot$} \\def\\msoly{$ M_\\odot~yr^{-1}$} \\def\\ltsima{$\\; \\buildrel < \\over \\sim \\;$} \\def\\simlt{\\lower.5ex\\hbox{\\ltsima}} \\def\\gtsima{$\\; \\buildrel > \\over \\sim \\;$} \\def\\rfe{R$_{\\rm FeII}$} \\def\\ai{A.I.($\\frac{1}{4}$)\\/} \\def\\aox{$\\alpha_{ox}$\\/} \\def\\gs{$\\Gamma_{\\rm soft}$} \\def\\gh{$\\Gamma_{\\rm hard}$} \\def\\G{$\\Gamma$} \\def\\ne{n$_{\\rm e}$\\/} \\def\\ergss{ergs s$^{-1}$} \\def\\simgt{\\lower.5ex\\hbox{\\gtsima}}            % \\def\\bv{{\\bf $\\bigvee$}\\/} \\def\\qu{{$\\lambda$ 1400}} \\def\\h{$h^{-1}$} \\def\\hii{{\\sc Hii\\/}} \\def\\hi{{\\sc Hi\\/}} \\def\\a{$\\alpha$} \\def\\ha{{\\sc H}$\\alpha$} \\def\\nv{{\\sc Nv}$\\lambda$1240} \\def\\civ{{\\sc{Civ}}$\\lambda$1549\\/} \\def\\civnc{{\\sc{Civ}}$\\lambda$1549$_{\\rm NC}$\\/} \\def\\civbc{{\\sc{Civ}}$\\lambda$1549$_{\\rm BC}$\\/} \\def\\cmsq{cm$^{-2}$\\/} \\def\\cm3{cm$^{-3}$\\/} \\def\\md{$\\dot{\\rm M}$} \\def\\rg{$\\rm R_{\\rm g}$} \\def\\hb{{\\sc{H}}$\\beta$\\/} \\def\\hg{{\\sc{H}}$\\gamma$\\/} \\def\\hbbc{{\\sc{H}}$\\beta_{\\rm BC}$\\/} \\begin{document} ",
        "introduction": "}  The Broad Line Region (BLR) in Active Galactic Nuclei (AGN) is unresolved with present day imaging detectors and it will remain so for the foreseeable future. This is why ``one quasar spectrum is really worth a thousand images'' as stressed  by Gary Ferland at this meeting. In response we would add that {\\em a thousand spectra  are better than one average spectrum}. Understanding the diversity in optical spectroscopic properties of AGN  is the key to any realistic AGN modeling  \\cite{bg92,smd}.  The ideal to reconstruct the BLR velocity field from a single profile is not realistic \\cite{pen90}.  Determination of BLR structure and kinematics can be approached in two ways. It has been recognized for a long time that strong broad and narrow emission lines coming from both high and low ionization species are present in Seyfert galaxies and quasars. This is considered a defining spectroscopic property of AGN.  Restricting attention to broad lines: a) typical (i.e., strongest and most frequently observed) high ionization lines (HIL: ionization potential $\\simgt$ 50 eV) are \\civ\\ \\heii\\ and \\heiiuv\\ lines; while b) observed low ionization lines (LIL:  ionization potential $\\simlt $ 20 eV) include \\hi\\ Balmer lines, \\feii\\ multiplets, \\mgii, and the Ca{\\sc ii}\\ IR triplet.  The  first approach involves the study of line variability in response to continuum changes. This approach  has been pursued through a number of successful monitoring campaigns using Reverberation Mapping (RM) techniques.  RM requires a large amount of telescope time and, consequently, has been achieved only for a handful of sources. RM confirms that photoionization is the main heating process in the BLR and that a large part of the BLR is optically thick to the ionizing continuum (e. g. \\pcite{baldwin97}). RM studies have quantified a main difference between HIL and LIL; HIL respond to continuum changes with a time delay of a few days while the LIL respond with a delay of tens of days. This implies  that the LIL are emitted at larger distance from the continuum source (\\pcite{ulr,kor95}.  An exhaustive list of reference can be found in \\pcite{smd}). RM applied to line profiles suffers from uncertainty in our knowledge about the physics of continuum and broad line formation, so that conflicting models can still describe the same lag times \\cite{pen91,wanpet96}.  The second approach involves statistical analysis of large samples of line profiles which differ because of properties that may affect the BLR structure,  for example samples of  radio quiet (RQ) and radio loud (RL) AGN. Statistical studies can be done for a single line or by comparing lines sensible to different physical parameters (e.g., strongest LIL and HIL). The statistical  approach, on which we will focus here, is more empirical and therefore requires a conceptual framework for interpretation. This approach is often criticized as relying on several assumptions including that the profile variability does not influence profile shapes and that the non-simultaneity of the observations of LIL and HIL are unimportant. Actually as the size of high quality data samples grow these effects become less and less important. ",
        "conclusions": "While observations support emission from an accretion disk and an associated spherical wind in RQ AGN with FWHM(\\hbbc) $\\simlt$ 4000 \\kms, there is not enough observational support to warrant the same conclusion for RL AGN (and possibly RQ with FWHM$>$4000 \\kms), although a disk + wind model is a viable possibility also in this case. Wide separation double peakers (mostly RL) do not provide conclusive evidence in favor of LIL disk emission; rather, there is evidence against disk emission as well as against every other reasonably simple scenario."
    },
    "0002/astro-ph0002433_arXiv.txt": {
        "abstract": "Based on the refined dynamical model proposed by us earlier for beamed $\\gamma$-ray burst ejecta, we carry out detailed numerical procedure to study those $\\gamma$-ray bursts with rapidly fading afterglows (i.e., $\\sim t^{-2}$). It is found that optical afterglows from GRB 970228, 980326, 980519, 990123, 990510 and 991208 can be satisfactorily fitted if the $\\gamma$-ray burst ejecta are highly collimated, with a universal initial half opening angle $\\theta_0 \\sim 0.1$. The obvious light curve break observed in GRB 990123 is due to the relativistic-Newtonian transition of the beamed ejecta, and the rapidly fading optical afterglows come from synchrotron emissions during the mildly relativistic and non-relativistic phases. We strongly suggest that the rapid fading of afterglows currently observed in some $\\gamma$-ray bursts is evidence for beaming in these cases. ",
        "introduction": "The cosmological origin of $\\gamma$-ray bursts (GRBs) has been well established due to recent discovery of multi-wavelength afterglows (Costa et al. 1997; Metzger et al. 1997; Galama et al. 1997; Wijers, Rees \\& M\\'{e}sz\\'{a}ros 1997; Piran 1999). However, we are still far from resolving the puzzle of GRBs, because their ``inner engines'' are well hidden from direct afterglow observations. Some GRBs localized by BeppoSAX satellite have implied isotropic energy release of more than $10^{54}$ ergs (Kulkarni et al. 1998, 1999; Andersen et al. 1999; Harrison et al. 1999), which forces many theorists to deduce that GRB radiation must be highly collimated. Obviously, whether GRBs are beamed or not has become one of the most important problems that need to be solved urgently.  In the literature, it is generally believed that afterglows from jetted GRB remnant are characterized by an obvious break in the light curve during the {\\em relativistic phase}, due to both the jet edge effect (Panaitescu \\& M\\'{e}sz\\'{a}ros 1999; Kulkarni et al. 1999; M\\'{e}sz\\'{a}ros \\& Rees 1999) and the lateral expansion effect (Rhoads 1997, 1999). The breaking point is determined by $\\gamma \\sim 1/\\theta$, where $\\gamma$ is the Lorentz factor of the jet and $\\theta$ is the half opening angle. Recently we have developed a refined dynamical model that can correctly describe the overall evolution of an ultra-relativistic jet to non-relativisitic phase with the expanding velocity as small as $10^{-3} c$ (Huang et al. 2000a). Surprisingly enough, our detailed numerical results (Huang, Dai \\& Lu 2000b) show that the break theoretically predicted in light curve does not appear during {\\em the relativistic phase}, i.e., the time determined by $\\gamma \\sim 1/\\theta$ is not a breaking point. However, an obvious break does appear within the relativistic-Newtonian transition region, the degree of which is found to be parameter dependent (Huang, Dai \\& Lu 2000b). Generally speaking, the Newtonian phase of jet evolution is characterized by a rapid decay of optical afterglows, with the power-law timing index $\\alpha \\sim 1.8$ --- 2.1.  In practical observations, the power-law decay indices of afterglows from GRB 980326, 980519 and 991208 are anomalously large, $\\alpha \\sim 2.0$ (Groot et al. 1998; Owens et al. 1998; Castro-Tirado et al. 1999b), and optical light curves of GRB 990123 and 990510 even show obvious steepening at observing time $t \\geq 1$ --- 2 d (Kulkarni et al. 1999;  Harrison et al. 1999; Castro-Tirado et al. 1999a). Recently GRB 970228 was also reported to have a large index of $\\alpha \\sim 1.73$ (Galama et al. 1999b). These phenomena have been widely regarded as evidence for beaming (Sari, Piran \\& Halpern 1999; Castro-Tirado et al. 1999a). The purpose of this {\\em Letter} is to study these cases numerically, based on our refined beaming model (Huang et al. 2000a). It is found that optical afterglows from these GRBs can be easily reproduced, thus a jet model is strongly favored.  \\begin{figure} \\begin{picture}(100,160) \\put(0,0){\\special{psfile=fig01.ps angle=-90 hoffset=-30 voffset=215 hscale=43 vscale=43 }} \\end{picture} \\caption { Observed optical (R band) afterglow from GRB 990123 (Galama et al. 1999a; Kulkarni et al. 1999; Fruchter et al. 1999; GCN 354). The solid line is our best fit to it by employing a jet model. Inset shows the evolution of $\\gamma$ in our model } \\end{figure} ",
        "conclusions": "Determining the existence of beaming in GRBs will be helpful to our understanding of the GRB ``central engines''. Based on our refined dynamical model for beamed GRB remnants (Huang et al. 2000a), we have examined those GRBs with rapidly decaying optical afterglows closely. Detailed numerical results show that afterglows from GRB 970228, 980326, 980519, 990123, 990510 and 991208 can be satisfactorily fitted: the obvious break in the optical light curve of GRB 990123 is due to the relativistic-Newtonian transition of the beamed ejecta, and the rapid fading of afterglows from other GRBs is due to the relatively large values of $\\xi_{\\rm e}, \\xi_{\\rm B}^2$ and $n$. We thus strongly suggest that these GRBs be highly collimated. Note that in all cases, synchrotron radiation during the mildly relativistic and non-relativistic phases plays an important role in explaining the rapid decaying of optical afterglows.  In our calculations, the $\\xi_{\\rm e}$ values distribute in a narrow range, i.e., between 0.1 and 0.2. This has given some support to Freedman \\& Waxman's (1999) suggestion that $\\xi_{\\rm e}$ has a universal value of $\\sim 1/3$. However our results do not support their proposal that $p$ has a universal value of 2.2. Our $\\theta_0$ values distributed in a narrow range, $\\theta_0 \\sim 0.1$. This may provide important clues to our understanding of the central engine. We note that Woosley et al. (1999) have obtained such a small value of $\\theta_0$ after numerically studying the collapsar models of GRBs. The range of our $n$ values ($n \\sim 1$ --- 1000 cm$^{-3}$) indicates that some GRBs may be in gaseous environments, also giving some hints on GRB central engines.  The jet model greatly relaxes the energy crisis for GRB 990123 and 990510. However, we should keep in mind that other GRBs such as GRB 970508, 971214, 980329 and 980703 do not have rapidly fading afterglows. They should not be highly collimated (Huang, Dai \\& Lu 2000b). Then the energy crisis is really a problem: GRB 971214 and 980703 have indicated isotropic $\\gamma$-ray energies of $\\sim 0.17$ M$_{\\odot} c^2$ and $\\sim 0.06$ M$_{\\odot} c^2$ respectively!  The rapid fading of afterglows from GRB 970228, 980326, 980519 has also been explained as being due to the interaction of an isotropic blastwave with a wind environment (Dai \\& Lu 1998; Chevalier \\& Li 1999). However the wind environment model could not explain the light curve break observed for GRB 990123 and 990510. As has been shown clearly in this paper, the jet model can naturally explain all these bursts, and should be more reasonable.  Another possibility was proposed recently by Dai \\& Lu (1999, 2000). They suggested that the light curve break is due to the relativistic-Newtonian transition of an isotropic blastwave in a dense medium ($n \\sim 10^6$ cm$^{-3}$), and the rapidly fading afterglows can be explained as emissions in the non-relativistic phase. It is obviously an interesting proposition and should also be paid attention to."
    },
    "0002/astro-ph0002427_arXiv.txt": {
        "abstract": "A modification of the popular unification model for Seyfert galaxies is proposed, which takes into account recent observational findings on the statistical properties of both type 1 and type 2 Seyferts.  In the proposed scenario, Compton--thick Seyfert 2 galaxies are those sources observed through compact, thick matter (the `torus'), while Compton--thin/ intermediate Seyfert galaxies are obscured by dust lanes at larger distances. ",
        "introduction": "The discovery of broad lines in the polarized flux of the archetypal Seyfert 2 galaxy NGC~1068 (Antonucci \\& Miller 1985) has been a landmark in the study of AGN, leading to the now widely accepted unification model for Seyfert galaxies (e.g. Antonucci 1993). In this scenario, type 1 and type 2 Seyferts are intrinsically the same, appearing different only because in the former we can see the nucleus directly, while in the latter the direct view is prevented by absorbing matter on the line--of--sight.  There can be no doubt about the basic correctness of the unification model. In fact, there is plenty of examples of Seyfert 2s with unambiguously have an obscured type 1 nucleus at their centre, while we are not aware of even a single Seyfert 2 which certainly does not harbour a hidden Seyfert 1. However, it is likely that the strictest version of the model (in which the aspect angle is the only relevant parameter) is not fully valid. Arguments against it include:  \\begin{itemize} \\item a) there is, on average, enhanced star formation in Seyfert 2 galaxies with respect to Seyfert 1s (Maiolino et al. 1997);  \\item b) the average morphologies between galaxies hosting type 1 and 2 nuclei are different, those hosting type 2 being on average more irregular (Maiolino et al. 1997, Malkan et al. 1998);  \\item c) there is a greater overall dust content in Seyfert 2s (Malkan et al. 1998). \\end{itemize}   \\begin{figure*}[t] \\centerline{\\epsfig{figure=Ca042.f1.ps,height=18.0cm, width=18.0cm}} \\caption{The proposed unification model (see text for details).} \\label{umplot} \\end{figure*}   Moreover, Malkan et al. (1998) showed that there is plenty of dust lanes at distances of hundred of parsecs in all type of Seyfert galaxies. They went as far as to propose that these dust lanes are completely responsible for the type 1/type 2 dichotomy, and therefore dismissing the existence of the torus. To avoid confusion, it is important to remark that here and after we use the term `torus' to indicate any distribution of optically thick matter close (a few tens of parsecs at most) to the nucleus, and with a large covering factor, whatever its actual geometry is (not necessarily ring--shaped!). Actually, there are many good arguments in favour of the existence of the torus. Apart from those listed by Antonucci (1993), more recent ones include radio imaging and water maser measurements, indicating dense matter very close to the black hole (e.g. Gallimore, Baum \\& O'Dea 1997; Greenhill et al. 1996); and infrared imaging of nearby Seyfert 2s,  again indicating the presence of large amount of matter very close to the nucleus (e.g. Siebenmorgen et al. 1997).  In the last few years, very strong evidence in favour of the `torus' (whatever it really is) has been obtained from X--rays observations. In particular, BeppoSAX observations have shown that at least half of Seyfert 2s in the local Universe are Compton--thick (Maiolino et al. 1998a; Risaliti, Maiolino \\& Salvati 1999), i.e. the nuclear radiation is absorbed by matter with column densities exceeding 10$^{24}$ cm$^{-2}$ (see Matt et al. 2000 for the general properties of bright Compton--thick Seyfert 2s). While in a handful of Compton--thick sources the column density has been directly measured (e.g. NGC~4945: Iwasawa et al. 1993, Done et al. 1996, Guainazzi et al. 2000; Circinus Galaxy: Matt et al. 1999; NGC~6240: Vignati et al. 1999), in the majority of them either the column density is so large to completely obscure the nucleus even in hard X--rays (e.g. NGC~1068: Matt et al. 1997) or their flux at high energies is too low to permit a detailed spectral analysis or, often, even a detection with the present generation of hard X--ray detectors. In this case, the classification of a source as Compton--thick lies on indirect arguments: a reflection--dominated spectrum (recognized by the flat slope, if the reflector is `cold', and by a $\\sim$1 keV equivalent width iron line) is the most useful and used indicator.   The observed large fraction of Compton--thick Seyfert 2s implies that the covering factor of such thick matter must be large. Assuming a spherical geometry for simplicity, the total amount of matter is proportional to the square of the outer radius, provided that it is much larger than the inner radius (this argument holds, at least roughly, whatever is the geometry, if the covering factor is large). In order not to exceed the value of the mass obtained from dynamical measurements, the outer radius of the torus in Circinus Galaxy must be less than 20 pc (Maiolino et al. 1998b). A less tight constraint is derived from NGC~1068 (Risaliti et al. 1999), i.e $\\simlt$100 pc, which however still implies that the dust lanes on the hundred parsecs scale  cannot be the matter responsible for the absorption in this source.  A further important finding of Risaliti et al. (1999) is that there is a clear difference between the $N_H$ distribution of Intermediate (type 1.8--1.9) and strict type 2 Seyferts. While the intermediate Seyferts in the Risaliti et al. sample are all Compton--thin, the strict type 2 Seyferts have column densities generally exceeding 10$^{23}$ cm$^{-2}$, and most of them are Compton--thick.  In the following section we will discuss a possible modification of the unification model which qualitatively accounts for the different statistical properties of obscured and unobscured Seyfert galaxies, and for the different column density distribution of intermediate and strict type 2 sources. ",
        "conclusions": "The proposed  modification, while retaining the basic and most successful characteristic of the unification model (i.e. that all intermediate and type 2 Seyferts harbour an obscured type 1 nucleus), explains, at least qualitatively, the observed differences between the average properties of Seyfert 1s and Seyfert 2s and between the column density distributions of intermediate and strict type 2 Seyferts. This model is somewhat different (even if on a similar line of thought) from that proposed by Maiolino \\& Rieke (1995), and takes into account recent observational results. The two `flavours' of Seyferts, with and without the thick torus, may represent either two different branches in the AGN evolution, or a different evolutionary stage in the life of any (or most) source.  Testing the proposed model would require further studies of the statistical properties of Seyferts, which caution in separating intermediate and strict type 2 sources. For instance, as the torus is expected to be axially symmetric (while the dust lanes are probably more randomly distributed), a correlation between Compton-thick absorption and presence of ionization cones and large polarization is expected. Another test is to search for the presence of the torus in Seyfert 1s, as we predict that many sources of this class do not have it.  This may be done either by searching for strong IR emission (which however may be related to starburst rather than reprocessing of UV/X--rays from thick matter), of by searching for signatures of X--ray reprocessing (e.g. Ghisellini, Haardt \\& Matt 1994; Krolik, Madau \\& \\.Zycki 1994). The latter way have already produced a clear case of a Seyfert 1 {\\it with} the torus: NGC~4051 was caught by BeppoSAX when the nuclear emission was switched--off, and clear evidence for reprocessing by large amount of optically thick distant matter was present (e.g. a $\\sim$600 eV equivalent width iron line and a cold reflection continuum, Guainazzi et al. 1998). However, this kind of observations based on variability requires rather extreme behaviours of the X--ray emission, and it is not clear how common a switching--off of the nucleus can be. More promising is to search for narrow (i.e. unresolved) iron K$\\alpha$ lines in addition to the relativistically broadened component from the accretion disc. When the nuclear X--ray emission is directly visible, equivalent widths of the order of 50--100 eV are expected from the torus  (Ghisellini, Haardt \\& Matt 1994). While ASCA and BeppoSAX results have been rather ambiguous in this respect, such a search is certainly within the capabilities of XMM."
    },
    "0002/astro-ph0002082_arXiv.txt": {
        "abstract": "Studies of the X-ray surface brightness profiles of clusters, coupled with theoretical considerations, suggest that the breaking of self-similarity in the hot gas results from an `entropy floor', established by some heating process, which affects the structure of the intracluster gas strongly in lower mass systems.  By fitting analytical models for the radial variation in gas density and temperature to X-ray spectral images from the \\emph{ROSAT} PSPC and \\emph{ASCA} GIS, we have derived gas entropy profiles for 20 galaxy clusters and groups. We show that when these profiles are scaled such that they should lie on top of one another in the case of self-similarity, the lowest mass systems have higher scaled entropy profiles than more massive systems.  This appears to be due to a baseline entropy of 70-140~${h_{50}}^{-\\frac{1}{3}}$ keV cm$^{2}$ depending on the extent to which shocks have been suppressed in low mass systems.  The extra entropy may be present in all systems, but is detectable only in poor clusters, where it is significant compared to the entropy generated by gravitational collapse. This excess entropy appears to be distributed uniformly with radius outside the central cooling regions.  We determine the energy associated with this entropy floor, by studying the net reduction in binding energy of the gas in low mass systems, and find that it corresponds to a preheating temperature of $\\sim 0.3$~keV. Since the relationship between entropy and energy injection depends upon gas density, we are able to combine the excesses of 70-140~keV~cm$^{2}$ and 0.3~keV to derive the typical electron density of the gas into which the energy was injected. The resulting value of 1-3$\\times 10^{-4}{h_{50}}^{\\frac{1}{2}}$\\,cm$^{-3}$, implies that the heating must have happened prior to cluster collapse but after a redshift $z\\sim$ 7-10.  The energy requirement is well matched to the energy from supernova explosions responsible for the metals which now pollute the intracluster gas. ",
        "introduction": "The hierarchical clustering model for the formation of structure in the universe predicts that dark matter halos should be scaled versions of each other\\cite{navarro95a}. While some energy transfer between dark matter and gas is possible through gravitational interaction and shock heating, simulations suggest that the gas and dark matter halos will be almost self-similar in the absence of additional heating or cooling processes \\cite{eke98a}. Comparison of the structure of real galaxy systems with this predicted self-similarity provides an excellent probe of extra physical processes that may be taking place in galaxy clusters and groups.  It has been suggested that specific energy in cluster cores is higher than expected from gravitational collapse and that this may be due to energy injected by supernova-driven protogalactic winds \\cite{white91a,david91a}. \\scite{david96a} studied the entropy in a small sample of galaxy systems and suggested that the entropy in their cores had been flattened due to energy injection.  \\scite{ponman98a} have recently shown that the surface brightness profiles of clusters and groups do not follow the predicted self-similar scaling.  Surface brightness profiles of galaxy groups are observed to be significantly flatter than those of clusters, indicating differences in the gas distribution.  \\begin{table*} \\begin{minipage}{154mm} \\caption{Some important properties for the sample of 20 galaxy clusters and groups. Systems are listed in order of increasing temperature.} \\label{tab:sample} \\begin{tabular}{lccccccc} \\hline \\hline Cluster/Group&R.A.(J2000)&Dec.(J2000)&z&$N_{H}$($10^{20}$ cm$^{-2}$)&$T$(keV)&$Z$(solar)&Data\\\\ \\hline HCG 68&208.420&40.319&0.0080&0.90&0.54&0.43&PSPC\\\\ HCG 97&356.845&-2.169&0.0218&3.29&0.87&0.12&PSPC\\\\ HCG 62&193.284&-9.224&0.0137&3.00&0.96&0.15&PSPC\\\\ NGC 5044 Group&198.595&-16.534&0.0082&5.00&0.98&0.27&PSPC\\\\ RX J0123.6+3315&20.921&33.261&0.0164&5.0&1.26&0.33&PSPC\\\\ Abell 262&28.191&36.157&0.0163&5.4&1.36&0.27&PSPC\\\\ IV Zw 038&16.868&32.462&0.0170&5.3&1.53&0.40&PSPC\\\\ Abell 400&44.412&6.006&0.0238&9.1&2.31&0.31&PSPC\\\\ Abell 1060&159.169&-27.521&0.0124&5.01&3.24&0.27&PSPC+GIS\\\\ MKW 3s&230.507&7.699&0.0453&3.1&3.68&0.30&PSPC\\\\ AWM 7&43.634&41.586&0.0173&9.19&3.75&0.33&PSPC\\\\ Abell 780&139.528&-12.099&0.0565&4.7&3.8&0.23&PSPC+GIS\\\\ Abell 2199&247.165&39.550&0.0299&0.87&4.10&0.30&PSPC\\\\ Abell 496&68.397&-13.246&0.0331&4.41&4.13&0.31&PSPC+GIS\\\\ Abell 1795&207.218&26.598&0.0622&1.16&5.88&0.26&PSPC\\\\ Abell 2218&248.970&66.214&0.1710&3.34&6.7&0.20&PSPC+GIS\\\\ Abell 478&63.359&10.466&0.0882&13.6&7.1&0.21&PSPC+GIS\\\\ Abell 665&127.739&65.854&0.1818&4.21&8.0&0.28&PSPC+GIS\\\\ Abell 1689&197.873&-1.336&0.1840&1.9&9.0&0.26&PSPC+GIS\\\\ Abell 2163&243.956&-6.150&0.2080&11.0&13.83&0.19&PSPC+GIS\\\\ \\hline \\end{tabular} Notes: Positions, hydrogen columns and redshifts are taken from \\scite{ebeling96a}, \\scite{ebeling98a}, \\scite{ponman96a} and \\scite{helsdon99a}. Emission weighted temperatures and metallicities for these systems are taken from \\scite{markevitch96a}, \\scite{mushotzky97a}, \\scite{markevitch98a}, \\scite{mchardy90a}, \\scite{fukazawa98a}, \\scite{butcher91a}, \\scite{mushotzky97b}, \\scite{helsdon99a}, \\scite{david96a}, \\scite{david94a} and \\scite{ponman96a}. \\end{minipage} \\end{table*}  In order to explore this effect further, it is necessary to study the properties of the gas in these systems in greater detail. A particularly interesting property of the gas for this purpose is its entropy, as this will be conserved during adiabatic collapse of the gas into a galaxy system, but is likely to be altered by any other physical processes. For instance preheating of the gas before it falls into the cluster, energy injection from galaxy winds and radiative cooling of the gas in dense cluster cores will all perturb the entropy profiles of clusters from the self-similar model. Analysis of the entropy profiles of virialized systems of different masses should therefore allow the magnitude of such effects to be studied, constraining the possible processes responsible. A key question to answer in this regard is how much energy is involved in any departures from self-similarity of the entropy profiles. The study of \\scite{ponman98a} was not able to address this issue in detail, since the gas was assumed to be isothermal.  Here we combine \\emph{ROSAT} PSPC and \\emph{ASCA} GIS data to constrain temperature profiles, allowing a more detailed study of entropy and energy distributions in the intergalactic medium (IGM).  Energy loss from the gas due to cooling flows in the centres of clusters and groups will actually lead to an increase in the gas entropy outside the cooling region \\cite{knight97a}. This is because as gas cools out at the centre of the system, gas from a larger radius, which has higher entropy, flows in adiabatically to replace it. However this effect will not be very large unless a significant fraction of the gas in the system cools out, which is not feasible within a Hubble time, even for systems with exceptionally large cooling flows.  Energy injection into the gas will also raise the entropy profiles of systems. This energy injection could occur either before or after the systems' collapse, but more energy is needed to get the same change in entropy when the gas is more dense \\cite{ponman98a}. There are several possible processes that might have injected energy into the intracluster medium: radiation from quasars, early population III stars, or energetic winds associated with galaxy formation.  There may also be transient effects on the entropy profiles of systems due to recent mergers. Hydrodynamical simulations suggest that the entropy profiles of systems are flattened and their central entropy raised during a merger, and this will last until the system settles back into equilibrium \\cite{metzler94a}. In order to look for the effects of extra physical processes, it is advantageous to study a set of systems with a large range in system mass, as these processes will break the expected self-similar scaling relations. In the present paper we examine the properties of the intracluster gas in systems with mean temperatures ranging over a factor of 25, corresponding to virial masses varying by over two orders of magnitude. ",
        "conclusions": "It is clear from Figures \\ref{fig:plot1} and \\ref{fig:plot2} that systems with integrated temperatures below 4 keV show signs of having excess entropy in their intracluster gas over what would be expected from the simple self-similar model. It can further be seen from Fig. \\ref{fig:plot3} that the amount of excess entropy does not depend systematically on the system temperature and, from Fig.  \\ref{fig:plot5}, it has an approximately constant value outside the central cooling flow regions.  The average excess entropy outside the cooling flow region lies in the range 70-140${h_{50}}^{-\\frac{1}{3}}$ keV cm$^{2}$. The upper limit, where shocks are totally suppressed in low mass systems, is comparable with the result of \\scite{ponman98a} who obtained a value of 100${h_{100}}^{-\\frac{1}{3}}$ (126${h_{50}}^{-\\frac{1}{3}}$)~keV~cm$^{2}$ for the assumption of total shock suppression. This new upper limit on the entropy should be more reliable as it does not rely on the assumption of isothermality of the intracluster gas that \\scite{ponman98a} had to use. Our analysis also sets a lower bound on the entropy for the case where the shock heating is not affected.  It is also interesting to compare our measured value for the excess entropy against the value assumed in various theoretical models of entropy injection in galaxy systems. For instance \\scite{balogh98a} assume a constant entropy injection value of $\\sim $350 keV cm$^{2}$ in order to reproduce the steepening in the $L$-$T$ relation for galaxy groups. \\scite{tozzi99a} argue that to steepen $L$-$T$ at $\\sim$ 0.5-2 keV, entropy injection in the range 190-960 keV cm$^{2}$ is needed. Both these values are somewhat higher than our measured range, but considering the simplified nature of these models, the similarity is encouraging. It will be interesting to see whether more sophisticated models can match the group $L$-$T$ relation using the lower values of entropy we observe.  A number of models work on the assumption of some specific amount of energy injection into the gas. These can be compared with the amount of excess energy we observe to be present in galaxy systems. \\scite{cavaliere97a} and \\scite{cavaliere98a} assume that the gas in galaxy systems is preheated to a temperature of 0.5 keV which is comparable to our measured value. \\scite{wu99a} obtain energy input of $<0.1$~keV per particle from SN heating within most of their hierarchical merger model runs. However, it is not clear that this represents a hard limit, since these authors assumed that gas can only be heated to the escape velocity of their galaxy halos. \\scite{wu98a,wu99a} also find that an injected energy per particle of $\\sim$1-2 keV is required to reproduce the slope of the cluster $L$-$T$ relation \\cite{david93a}. This may indicate that the preheating required to match the steepening in $L$-$T$ at $T<1$~keV does not provide a solution to the departure of the cluster relation from the self-similar result, $L\\propto T^2$. For example, it is clear that the model of \\scite{cavaliere98a}, which provides a good match to the group data, fails to reproduce the slope of the $L$-$T$ relation at high temperatures (see their Fig.9). Additional effects may be at work -- for example \\scite{allen98a} have demonstrated that allowing for the impact of cooling flows flattens the $L$-$T$ relation for rich clusters towards $L\\propto T^2$.  The floor entropy of 70-140 ${h_{50}}^{-\\frac{1}{3}}$~keV is small compared to the entropy of the 8 systems with temperatures of 4 keV or above, which averages 380 keV cm$^{2}$ (at 0.1 $R_{v}$). Hence our results are consistent with the idea that an approximately constant amount of excess entropy, $\\sim$70-140${h_{50}}^{-\\frac{1}{3}}$ keV cm$^{2}$, is present in {\\it all} of the systems, but is only noticeable in systems where it constitutes a large fraction of the total entropy, i.e. in systems with temperatures below 4 keV. From Fig. \\ref{fig:plot5} it can be seen that there is little evidence for any dependence of the excess entropy on radius outside the central cooling region.  This suggests that the process involved in injecting entropy into the systems does so fairly uniformly, at least within $0.25 R_{v}$.  From Fig. \\ref{fig:plot6}, the gas in low mass systems is significantly less tightly bound than would be expected from self-similar scaling. Combining the excess entropy and energy requirements leads us to conclude that the energy was injected at $z<$7-10, but before cluster collapse. Possible candidates for the source of this extra energy are preheating by quasars, population III stars or galaxy winds. It is known that since recombination at z $\\sim 1400$, the intergalactic medium has been re-ionized. This re-ionization is normally assumed to be caused by quasars or an early epoch of star formation. However analytical models of these processes \\cite{valageas99a,tegmark93a} suggest that the IGM will only be heated to $\\sim$ $10^{4}$-$10^{5}$K, resulting in an entropy change that is at least an order of magnitude lower than the measured value. In contrast, energy injected by supernovae associated with the formation of the bulk of galactic stars should be much more significant \\cite{white91a,david91a}.  The likely energies involved can be estimated from observed metal abundances in the intracluster gas. The major uncertainty here lies in establishing the contributions from supernovae of type Ia and type II, which have very different ratios of iron yield to energy \\cite{renzini93a}. Recent studies with \\emph{ASCA} \\cite{finoguenov99a,finoguenov99b} in which contributions from SNIa and SNII have been mapped in a sample of groups and clusters, by tracing the abundance of iron and alpha elements, leads to the conclusion that SNIa provide a significant contribution to the iron abundance, particularly in galaxy groups. The supernova energy associated with the observed metal masses by \\scite{finoguenov99a} are in good agreement with the energy of $\\sim 0.4$~keV per particle derived above on the basis of the observed energy excesses. This is also similar to the preheating involved in the models of \\scite{cavaliere97a}, \\scite{cavaliere98a} and \\scite{balogh98a} supporting the idea that the similarity breaking we see in the intracluster gas does result from preheating associated with galaxy formation.  With the forthcoming availability of data from \\emph{Chandra} and \\emph{XMM}, much more detailed studies of the abundance and entropy distributions of galaxy systems will become possible. This will allow deviations from mean trends to be studied in detail. Since galaxy winds will inject both energy and metals, whereas processes such as ram pressure stripping will lead to metal enrichment without heating, studies with these new X-ray observatories should throw a great deal of light on the evolutionary history of galactic systems and the galaxies they contain."
    },
    "0002/astro-ph0002031_arXiv.txt": {
        "abstract": "We review the relevance of ${\\rm H_2}$ molecules for structure formation in cosmology.  Molecules are important at high--redshifts, when the first collapsed structures appear with typical temperatures of a few hundred Kelvin. In these chemically pristine clouds, radiative cooling is dominated by ${\\rm H_2}$ molecules. As a result, ${\\rm H_2}$ ``astro--chemistry'' is likely to determine the epoch when the first astrophysical objects appear. We summarize results of recent three--dimensional simulations. A discussion of the effects of feedback, and implications for the reionization of the universe is also given. ",
        "introduction": "In current ``best--fit'' cosmological models, cold dark matter (CDM) dominates the dynamics of structure formation, and processes the initial density fluctuation power spectrum $P(K)\\propto k^n$ with $n=1$ to predict $n=1$ on large scales and $n\\approx-3$ on small scales (Peebles 1982).  The r.m.s. density fluctuation $\\sigma_M$ then varies inversely with the mass--scale ($\\sigma_M\\propto M^{-2/3}$ for $M\\gg 10^{12}\\msun$, while the dependence is only logarithmic for $M\\ll 10^{12}\\msun$).  The more overdense a region, the earlier it collapses, implying that the present structure was built from the bottom up, with smaller objects appearing first, and subsequently merging and/or clustering together to assemble the larger objects (Peebles 1980).  The predicted formation epochs of ``objects'' (i.e. collapsed dark matter halos) with various masses in the so--called standard CDM cosmology (Bardeen et al. 1986) are shown in Figure~\\ref{fig:zcollintro}.  Galaxies, which have masses around $10^{11-12}\\msun$, are expected to have formed when the universe had approximately 10\\% of its present age (redshift $z\\sim3$), just around the limit of the deepest present--day observations (i.e. the Hubble Deep Field, HDF, Williams et al. 1996; or Ly$\\alpha$ emission line detections, Weymann et al. 1998).  Clusters of galaxies with masses around $10^{14-15}\\msun$ are predicted to have formed as recently as 80\\% of the current age, with the more massive clusters still assembling at the present time.  \\begin{figure} \\centerline{\\psfig{figure=zcoll.ps,width=10cm,angle=0}} \\caption{\\label{fig:zcollintro} The formation epoch of objects with various masses in a standard CDM model ($\\Omega=n=\\sigma_{8h^{-1}}=1$), shown as a fraction of the current age of the universe.  The upper labels indicate redshifts. The first objects form at redshift $z\\sim 20$, at a ``lookback'' fraction of $\\sim$99\\%.} \\end{figure}  Objects with the masses of globular clusters, $10^{5-6}\\msun$, are predicted to have condensed as early as $\\sim$1\\% of the current age, or $z\\sim 20$.  It is natural to identify these condensations as the sites where the first ``astrophysical'' objects (stars, or quasars) might be born.  Although the CDM model in Figure~\\ref{fig:zcollintro} predicts still smaller condensations at even earlier times, the cosmological Jeans mass in the smooth gas after recombination is $\\sim 10^{4}\\msun$ (Peebles 1965), implying that gas pressure inhibits the collapse of gas below this scale (see Haiman, Thoul \\& Loeb 1996 on collapse on somewhat smaller scales due to gas/DM shell--crossing, hereafter HTL96).  What happens in a newly collapsed halo?  Formally, in the absence of non--gravitational forces, a perfectly spherical top--hat perturbation simply collapses to a point (Peebles 1980).  According to more accurate treatments describing self--similar solutions of spherical, but inhomogeneous secondary infall for a mixture of cold dark matter and baryons (Gunn \\& Gott 1972, Fillmore \\& Goldreich 1984, Bertschinger 1985), the evolution is as follows. In the initial stages, while the bulk of the cloud is still turning around or expanding, the central, densest region of the cloud collapses and forms a dense virialized clump.  The dark matter component virializes through violent relaxation (Lynden--Bell 1967), while the kinetic energy of the gas is converted into thermal energy through a hydrodynamic shock that raises the gas temperature to its virial value.  As subsequent gas shells fall towards the center, they encounter an outward--propagating shock, and are brought to a sudden a halt. Continuous accretion onto the center then establishes a stationary, virialized object with a $\\rho\\propto r^{-2.25}$ density profile.  Early discussions of the formation of galaxies and clusters have argued that the subsequent behavior of the gas in such a virialized object is determined by its ability to cool radiatively on a dynamical time (Rees \\& Ostriker 1977; Dekel \\& Silk 1986). The same ideas apply on the smaller scales expected for the first collapsed clouds (Silk 1977; Kashlinsky \\& Rees 1983).  Objects that are unable to cool and radiate away their thermal energy maintain their pressure support and identity, until they become part of a larger object via accretion or mergers.  On the other hand, objects that can radiate efficiently will cool and continue collapsing.  The cooling time is determined by the virial temperature, $T_{\\rm vir} \\sim 10^4{\\rm K}$ $(M/10^8{\\rm M_\\odot})^{2/3}$ $[(1+z)/11]$.  For the largest halos, above $T_{\\rm vir}\\ga10^6$K, the most important cooling mechanism is Bremsstrahlung; for galaxy--sized halos ($10^4$K$\\la T_{\\rm vir}\\la10^6$K), cooling is possible via collisional excitation of neutral H and He. As first pointed out by Saslaw \\& Zipoy (1967) and Peebles and Dicke (1968), the virial temperatures of the first clouds are below $10^4$K. In the primordial gas, the only molecule that could have a sufficient abundance is ${\\rm H_2}$, allowing cooling between $10^2$K$\\la T\\la10^4$K (see Dalgarno \\& Lepp 1987 for a review of astrochemistry in the early universe, and Stancil, Lepp and Dalgarno 1996 on the possible importance of other molecules, such as HD and LiH). Below $T\\la10^2$K, the collapsed clouds are unable to cool, and remain pressure--supported for longer than a Hubble time.  Although ${\\rm H_2}$ molecules are unimportant in the formation of structures on galactic scales, they likely play a key role in the formation of the first, smaller structures.  In particular, the abundance of ${\\rm H_2}$ controls the minimum sizes and formation times of the very first systems (HTL96, Tegmark et al. 1997).  The main issues regarding ${\\rm H_2}$ molecules in structure formation, addressed in this review, are:  \\begin{itemize}  \\item What is the ${\\rm H_2}$ abundance in the first collapsed objects?  \\item What is the parameter space (in mass, redshift and metallicity) when ${\\rm H_2}$ cooling dominates over all other cooling mechanisms?  \\item What is the parameter space when the ${\\rm H_2}$ cooling time is shorter than the dynamical time, so that ${\\rm H_2}$ can effect the dynamics of a system?  \\item What feedback mechanisms effect the ${\\rm H_2}$ abundance once the first stars or quasars light up?   \\end{itemize} ",
        "conclusions": "In popular CDM models, the first stars or quasars likely appeared inside condensed clumps with virial temperatures $T{\\rm vir}\\la 10^4$K at redshifts $z\\sim 20$.  Because of the low virial temperatures, ${\\rm H_2}$ cooling (or lack thereof) played a dominant role in the gas dynamics inside these condensations.  State--of--the art numerical simulations identify the sites for the first star--formation with the intersection of dense filaments (Abel et al. 1998, 2000).  Although not directly visible in ${\\rm H_2}$ emission, these star--formation sites could be detected out to $z\\sim 15$ with the Next Generation Space Telescope, provided they have a star formation efficiency of $\\ga$ one percent (Haiman \\& Loeb 1998), or they form quasar black holes with an efficiency of $\\sim 10^{-3}$ and shine at the Eddington luminosity (Haiman \\& Loeb 1998).  In the latter case, early mini--quasars from ${\\rm H_2}$ cooling would reionize intergalactic hydrogen by $z\\sim 10$.  The average star--formation efficiency in collapsed halos before $z\\sim 3$ can be estimated by matching the average metal enrichment of the Ly $\\alpha$ forest.  If the latter is $\\sim 10^{-3}$ solar, this implies that $\\sim$2\\% of the gas mass in collapsed regions are processed through stars (assuming a ratio in the number of high--mass to intermediate--mass stars as in a Scalo IMF, see Haiman \\& Loeb 1997).  This value is not far from the efficiency of $\\sim$1\\% suggested by 3D simulations that resolve sub--parsec scales (ABN00); however it is unlikely that the early IMF was similar to the one observed in the local universe (e.g. Larson 1999).  Arguably, further progress towards answering this question will have to come from a combination of more accurate simulations of star--formation in a metal--free plasma; including realistic radiative transfer."
    },
    "0002/astro-ph0002207_arXiv.txt": {
        "abstract": "We present the main results of a study of time-resolved spectra of 43 intense GRBs detected by BATSE. We considered the 4-parameter Band model and the Optically Thin Synchrotron Shock model (OTSSM). We find that the large majority of time-resolved spectra of GRBs are in remarkable agreement with the OTSSM. However, about  15 \\% of {\\it  initial  GRB pulses} show an apparent low-energy photon suppression. This phenomenon indicates that complex radiative conditions modifying optically thin emission may occur during the initial phases of some GRBs. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002494_arXiv.txt": {
        "abstract": "Near-IR spectra for a sample of 31 elliptical galaxies in the Coma cluster are obtained. The galaxies are selected to be ellipticals (no lenticulars), with a large spatial distribution, covering both the core and outskirt of the cluster (ie. corresponding to regions with large density contrasts). CO$_{sp}$ (2.3 $\\mu$m) absorption indices], measuring the contribution from intermediate-age red giant and supergiant stars to the near-IR light of the ellipticals, are then estimated.  It is found that the strength of CO$_{sp}$ features in elliptical galaxies increases from the core ($r < 0.2^\\circ$) to the outskirts ($r > 0.2^\\circ$) of the Coma cluster. Using the Mg$_2$ strengths, it is shown that the observed effect is not due to metallicity and is mostly caused by the presence of a younger population (giant and supergiant stars) in ellipticals in outskirts (low density region) of the cluster.  Using the CO$_{sp}$ features, the origin of the scatter on the near-IR Fundamental Plane (relation between the effective diameter, effective surface brightness and velocity dispersion) of elliptical galaxies is studied. Correcting this relation for contributions from the red giant and supergiant stars, the {\\it rms} scatter reduces from 0.077dex to 0.073dex. Although measurable, the contribution from these intermediate-age stars to the scatter on the near-IR Fundamental Plane of ellipticals is only marginal.  A relation is found between the CO$_{sp}$ and $V-K$ colours of ellipticals, corresponding to a slope of $0.036\\pm 0.016$, significantly shallower than that from the Mg$_2-(V-K)$ relation. This is studied using stellar synthesis models. ",
        "introduction": " ",
        "conclusions": "The  CO$_{sp}$ (2.3 $\\mu$m) absorption features are estimated from the near-IR spectra of a sample of 31 elliptical galaxies at the core and outskirts of the Coma cluster. Combined with other spectroscopic ($\\sigma$ and Mg$_2$) and photometric ( K-band) data for this sample, a study of the stellar population in elliptical galaxies is carried out. The main conclusions from this study are summarised as follows: \\begin{enumerate} \\item The mean CO$_{sp}$ values for elliptical galaxies at the core of the Coma cluster are found to be smaller compared to their counterparts in the outer region. There is a probability of $< 1\\%$ for these galaxies belonging to the same parent population. This is interpreted as due to the presence of intermediate-age red giant and supergiant stars in ellipticals in low density environments. This implies that the ellipticals in the outskirts of rich clusters are relatively younger than their counterparts at the core.  \\item Using the  CO$_{sp}$ values, the near-IR FP of ellipticals is corrected for contributions from the intermediate-age supergiant stars. This reduces the {\\it rms} scatter in this relation (at a given $log(A_e)$) from 0.077 dex to 0.073 dex. This modest reduction means that the observed scatter on the near-IR FP of ellipticals is mostly dominated by observational errors or dynamical effects and non-homology of ellipticals.  \\item A relation is found between CO$_{sp}$ and $V-K$ colours of ellipticals in the Coma cluster. The slope of this relation ($0.036\\pm 0.016$) is significantly shallower than that for the Mg$_2-(V-K)$ relation. Comparing with the stellar population models, it is found that the observed trend is caused by either metallicity or stellar population changes.  \\end{enumerate}"
    },
    "0002/astro-ph0002388_arXiv.txt": {
        "abstract": "The number density and clustering properties of Lyman-break galaxies (LBGs) observed at redshift $z\\sim 3$ are best explained by assuming that they are associated with the most massive haloes at $z\\sim 3$ predicted in hierarchical models of structure formation. In this paper we study, under the same assumption, how star formation and chemical enrichment may have proceeded in the LBG population. A consistent model, in which the amount of cold gas available for star formation must be regulated, is suggested. It is found that gas cooling in dark haloes provides a natural regulation process. In this model, the star formation rate in an LBG host halo is roughly constant over about 1 Gyr. The predicted star formation rates and effective radii are consistent with observations. The metallicity of the gas associated with an LBG is roughly equal to the chemical yield, or about the order of $1 Z_{\\odot}$ for a Salpeter IMF. The contribution to the total metals of LBGs is roughly consistent with that obtained from the observed cosmic star formation history. The model predicts a marked radial metallicity gradient in a galaxy, with the gas in the outer region having much lower metallicity. As a result, the metallicities for the damped Lyman-alpha absorption systems expected from the LBG population are low. Since LBG halos are filled with hot gas in this model, their contributions to the soft X-ray background and to the UV ionization background are calculated and discussed. ",
        "introduction": " ",
        "conclusions": "In this paper, we have examined the star formation and chemical enrichment in Lyman break galaxies, assuming them to be the central galaxies of massive haloes at $z\\sim 3$ and using simple chemical evolution models. We found that gas cooling in dark haloes provides a natural process which regulates the amount of star forming gas. The predicted star formation rates and effective radii are consistent with observations. The metallicity of the gas associated with an LBG is roughly equal to the chemical yield, or about the order of $1 Z_{\\odot}$ for a Salpeter IMF. Because of the relatively long star-formation time, the colours of these galaxies should be redder than that of short starbursts. It is not clear whether this prediction is consistent with current (rather) limited observations, because the interpretation of the observational data depends strongly on the adopted dust reddening. Stringent constraint can be obtained when full spectral information of the LBG population is carefully analyzed.  The model predicts a marked radial metallicity gradient in an LBG, with the gas in the outer region having lower metallicity. As a result, the metallicities for the damped Lyman-alpha absorption systems expected from the LBG population are lower than those for the LBGs themselves, although high metallicity is expected for a small number of sightlines going through the central regions of an LBG. At the same time, our modeled contribution to the total metal is roughly consistent with that obtained from the observed cosmic star formation history, i.e., there might not exist so-called ``missing metal\" problem although there could be more than half of the metals to be in the hot phase. Finally, a prediction of our model is that LBG haloes are filled with hot gas. As a result, these galaxies may have a non-negligible contribution to the soft X-ray background. The contribution of LBGs to the ionizing UV background is found to be small.  There are two basic assumptions in our work. One is that the LBG population is one-to-one associated with the most massive halos which are generated from the PS formalism, as done by MMWb; another is that the timescale of star formation for LBG population is assumed to be the order of 1Gyr, which is suggested by Steidel et al. (1999a,b, 1995). However, Baugh et al. (1999) recently argue that the prediction of the clustering properties of LBGs based on this first simple assumption will be discrepancy with the results of more detailed semi-analytic models.  Still, the second will lead to difficulty in reproducing the redshift evolution of bright galaxies (Kolatt et al. 1999). More detailed modelling done by Somerville (1997) suggest the collisional starbursts could be expected to be an important effect in understanding the LBGs. So, further observations are required to investigate the intrinsic properties of LBGs."
    },
    "0002/astro-ph0002341_arXiv.txt": {
        "abstract": "We present the results of a search for variable stars in the faint sparse globular cluster E3. We have found two variable stars: an SX Phe variable (V1) and a W UMa eclipsing binary (V2).  We have applied period-luminosity and period-color-luminosity relations to the variables to obtain their distance moduli. V1 seems to be a blue straggler belonging to E3, based on its distance modulus and location on the CMD. V2 is probably located behind the cluster, in the Milky Way halo.  We also present $V/B-V$ and $V/V-I$ color magnitude diagrams of E3. ",
        "introduction": "We present the results of a search for variable stars in the faint sparse globular cluster E3, located at $\\alpha_{2000}=9^h20^m59^s$, $\\delta_{2000}=-77\\arcdeg16\\arcmin57\\arcsec$. The cluster was discovered on the ESO B Schmidt Survey of the Southern Sky by Lauberts (1976). The first $BV$ photometry of the cluster was presented by van den Bergh, Demers \\& Kunkel (1980).  Numerous candidates for blue stragglers were identified. The photoelectric photometry of Frogel \\& Twarog (1983) confirmed this finding. A subsequent study by Hesser et al. (1984) using $UBV$ photoelectric and photographic observations showed a sparsely populated subgiant branch in the color-magnitude diagram.  The first CCD photometry for E3 in the $BV$ bands, obtained with the CTIO 4m telescope, was published by McClure et al. (1985). These observations suggested the presence of a second sequence of stars $\\sim0.75$ mag. above the cluster main sequence, interpreted as evidence for a significant population of binary stars in E3. The cluster was further studied by Gratton \\& Ortolani (1987), who provided new $BV$ CCD photometry from the 2.2m telescope at ESO.  \\begin{figure}[t] \\plotfiddle{mochejska.fig1.ps}{10cm}{0}{80}{80}{-180}{-10} \\caption{Finding chart for the variables in E3. North is up and east is to the left.} \\label{fig:map} \\end{figure} ",
        "conclusions": "Our variability search in E3 resulted in the discovery of two variable stars: an SX Phe variable (V1) and a W UMa eclipsing binary (V2).  We have applied period-luminosity and period-color-luminosity relations to the variables to obtain their distance moduli. V1 seems to be a blue straggler belonging to E3, based on its distance modulus and location on the CMD. V2 is probably located behind the cluster."
    },
    "0002/astro-ph0002306_arXiv.txt": {
        "abstract": "The blast-wave model for gamma-ray bursts (GRBs) has been called into question by observations of spectra from GRBs that are harder than can be produced through optically thin synchrotron emission.  If GRBs originate from the collapse of massive stars, then circumstellar clouds near burst sources will be illuminated by intense $\\gamma$ radiation, and the electrons in these clouds will be rapidly scattered to energies as large as several hundred keV.  Low-energy photons that subsequently pass through the hot plasma will be scattered to higher energies, thus hardening the intrisic spectrum.  This effect resolves the ``line-of-death'' objection to the synchrotron shock model. Illuminated clouds near GRBs will form relativistic plasmas containing large numbers of electron-positron pairs that can be detected within $\\sim$ 1-2 days of the explosion before expanding and dissipating.  Localized regions of pair annihilation radiation in the Galaxy would reveal past GRB explosions. ",
        "introduction": "The identification of flaring and fading X-ray, optical and radio counterparts to gamma-ray burst (GRB) sources (e.g., Costa et al.\\ \\markcite{cea97}1997; van Paradijs et al.\\ \\markcite{vpea97}1997; Djorgovski et al.\\ \\markcite{dea97}1997; Frail et al.\\ \\markcite{fea97}1997), and the large energy releases implied by redshift measurements, find a consistent explanation in an expanding relativistic blast-wave model (Paczy\\'nski \\& Rhoads \\markcite{pr93}1993; M\\'esz\\'aros \\& Rees \\markcite{mr97}1997).  As a result of Beppo-SAX and optical follow-on observations, the redshifts of about one dozen GRBs with durations greater than $\\sim 1$ s have been measured.  The distribution of redshifts is broad and centered near $z\\sim 1$, corresponding to the cosmological epoch of active star formation (Hogg \\& Fruchter \\markcite{hf99}1999).  GRBs are extremely luminous and energetic at hard X-ray and $\\gamma$-ray energies.  The degree of GRB collimation is unknown, but peak directional $\\gamma$-ray luminosities and energy releases as large as $\\partial L/\\partial \\Omega\\simeq 3\\times 10^{51}$ ergs (s-sr)$^{-1}$ and $\\partial E/\\partial \\Omega \\simeq 3\\times 10^{53}$ ergs sr$^{-1}$, respectively, have been measured (Kulkarni et al.\\ \\markcite{kea99}1999).  Less powerful GRBs and less luminous episodes during the GRB produce smaller $\\gamma$-ray powers, but the apparent isotropic $\\gamma$-ray luminosities from typical GRBs could regularly reach values exceeding $10^{50}L_{50}$ ergs s$^{-1}$ with $L_{50} \\sim 1$, with some GRBs reaching $L_{50} > 10^2$.  Because the energy radiated in $\\gamma$ rays is less than the total energy released by a GRB, the apparent isotropic energy release of GRB sources could often reach values of $10^{54}E_{54}$ ergs, with $E_{54} \\sim$1.  Cosmological gamma-ray burst and afterglow observations are best explained through the fireball/blast-wave model, where the deposition of large quantities of energy into a small region yields a fireball that expands until it reaches a relativistic speed determined by the amount of baryons mixed into the fireball (see, e.g., Piran \\markcite{tp99}1999 for a review).  Nonthermal synchrotron radiation from energetic electrons in the relativistic blast wave is thought to account for the origin of the prompt $\\gamma$-ray emission and afterglow radiation.  This paradigm has been called into question, however, by observations of very hard X-ray emission during the prompt $\\gamma$-ray luminous phase of a significant number of GRBs (Crider et al.\\ \\markcite{crider97}1997; Preece et al.\\ \\markcite{pea98}1998).  Photon fluxes $\\phi(\\e ) \\propto \\e^{-\\alpha_X}$ with $\\alpha_X \\sim$ 0, where $\\e = h\\nu/m_ec^2$ is the dimensionless photon energy, have been observed in 5-10\\% of GRBs that are bright enough to permit spectral analysis.  This strongly contradicts the optically-thin synchrotron shock model, which predicts that only radiation spectra with $\\alpha_X \\geq 2/3$ can emerge from the blast wave. In view of the severity of this challenge to the model, these observations have been termed the ``line-of-death\" to the synchrotron shock model. Possible explanations for this phenomeonon involve photoelectric absorption by optically thick cold matter (Liang \\& Kargatis \\markcite{lk94}1994; Brainerd \\markcite{jb94}1994; B\\\"ottcher et al.\\ \\markcite{bea99}1999), synchrotron self-absorption (Crider \\& Liang \\markcite{cl99}1999, Granot, Piran, \\& Sari \\markcite{granot00}2000, Lloyd \\& Petrosian \\markcite{lp99}1999), Compton scattering (Liang \\markcite{liang97}1997; Liang et al.\\ \\markcite{lcbs99}1999), or the existence of a pair-photosphere (M\\'esz\\'aros \\& Rees \\markcite{mr00}2000) within the blast wave.  Except for the last model cited, these explanations are inconsistent with the standard synchrotron shock model.  Here we offer a solution to this problem that is consistent with the standard model and recent observations pointing to a massive star origin of GRBs. ",
        "conclusions": ""
    },
    "0002/astro-ph0002130_arXiv.txt": {
        "abstract": "This paper describes an investigation of the early evolution of extra-galactic radio sources using samples of faint and bright Gigahertz Peaked Spectrum (GPS) and Compact Steep Spectrum (CSS) radio galaxies. Correlations found between their peak frequency, peak flux density and angular size provide strong evidence that synchrotron self absorption is the cause of the spectral turnovers, and indicate that young radio sources evolve in a self-similar way. In addition, the data seem to suggest that the sources are in equipartition while they evolve. If GPS sources evolve to large size radio sources, their redshift dependent birth-functions should be the same. Therefore, since the lifetimes of radio sources are thought to be short compared to the Hubble time, the observed difference in redshift distribution between GPS and large size sources must be due to a difference in slope of their luminosity functions. We argue that this slope is strongly affected by the luminosity evolution of the individual sources. A scenario for the luminosity evolution is proposed in which GPS sources increase in luminosity and large scale radio sources decrease in luminosity with time. This evolution scenario is expected for a ram-pressure confined radio source in a surrounding medium with a King profile density. In the inner parts of the King profile, the density of the medium is constant and the radio source builds up its luminosity, but after it grows large enough the density of the surrounding medium declines and the luminosity of the radio source decreases. A comparison of the local luminosity function (LLF) of GPS galaxies with that of extended sources is a good test for this evolution scenario. Unfortunately, only a handful of GPS sources are known at low redshift, and an LLF can only be derived, assuming that their cosmological number density evolution is similar to that of steep spectrum sources. The LLF derived in this way is shown to be in good agreement with the proposed evolution scenario. However, the uncertainties are large, and larger, homogeneously selected samples of GPS sources are needed. ",
        "introduction": "\\subsection{Gigahertz Peaked Spectrum and Compact Symmetric Objects.}  An important element in the investigation of the evolution of extra-galactic radio sources is the study of young counterparts of `old' FRI/FRII extended objects. Two classes of compact radio source can be found in the literature as most likely representatives of this early evolutionary stage: I) Gigahertz Peaked Spectrum (GPS) sources, which are characterised by a convex-shaped radio spectrum peaking at about 1 GHz in frequency (O'Dea 1998), and II) Compact Symmetric Objects (CSO) which are characterised by their small size ($<500$ pc) and two-sided radio structure, e.g. having jets and lobes on both sides of a central core (Wilkinson et al. 1994). Clearly, samples of GPS sources and CSOs are selected in very different ways.  GPS sources are selected on their broadband radio spectra, while CSOs are selected on their multi-frequency milli-arcsecond morphology. Therefore studies of these objects have mostly been presented separately.  However, a significant overlap between the two classes of sources exists.  GPS sources optically identified with galaxies are most likely to possess compact symmetric radio morphologies (Stanghellini et al. 1997a, 1999), and the large majority of CSOs exhibit a gigahertz-peaked spectrum. The large but not complete overlap between these two classes of source is most likely caused by the synchrotron self-absorbed mini-lobes, located at the extremities of most CSOs, being the main contributors to the overall radio spectrum, and producing the peak at about 1 GHz in frequency.  \\subsection{Evidence for the young nature of GPS sources and CSOs.}  Since the early discovery of GPS sources, it has been speculated that these were young objects (Shklovsky 1965, Blake 1970). However, a commonly discussed alternative to them being young was that they are small due to confinement by a particularly dense and clumpy interstellar medium that impedes the outward propagation of the jets (van Breugel, Miley \\& Heckman 1984; O'Dea, Baum \\& Stanghellini 1991). This latter hypothesis now looks less likely since recent observations show that the surrounding media of peaked spectrum sources are not significantly different from large scale radio sources, and insufficiently dense to confine these sources. The most compelling evidence for youth is found in observations of the propagation velocities of the hot spots of several GPS/CSOs (Owsianik \\& Conway 1998; Owsianik, Conway \\& Polatidis 1998; Tschager et al. 1999). They all appear to have separation velocities of typically $\\sim 0.2h^{-1}c$, indicating a dynamical age of $\\sim 10^3$ year and clearly showing that these are indeed young objects. Recent measurements of the high frequency breaks in  Compact Steep Spectrum (CSS) sources, indicate that these somewhat larger objects have radiative ages in the range of $10^3$ to $10^5$ years (Murgia et al. 1999).   \\subsection{Current views on Radio Source Evolution}  Observational constraints on the luminosity evolution of radio sources mainly come from the source density in the power - linear size $(P-D)$ diagram (Shklovsky 1963). It was found that sources with large sizes ($D>1$ Mpc, eg. Schoenmakers 1999) and high radio luminosities ($P>10^{26}$ W/Hz at 178 MHz) are rare, suggesting that the luminosities of sources should decrease quickly with linear sizes approaching 1 Mpc. Several authors have compared the number densities of GPS and CSS sources  with those of large radio sources to investigate the luminosity evolution of young radio sources (Fanti et al. 1995; Readhead et al. 1996; O'Dea \\& Baum 1997). Fanti et al. (1995) argue that the luminosities of CSS sources decrease by a factor of $\\sim 10$ as they evolve to extended objects. Readhead et al. (1996) find a factor of 8 decrease in luminosity as a source expands from 500 pc to 200 kpc in overall size. Taking into account their CSO statistics, they find that the luminosity evolution in the phases CSO-MSO-LSO (MSO = Medium Symmetric Object; LSO = Large Symmetric Object), i.e. from 10 pc to 150 kpc, is consistent with a single power-law luminosity evolution. This conclusion is not supported by O'Dea \\& Baum (1997), who found that GPS and CSS sources must decrease in luminosity at a faster rate than the classical 3CR doubles. Blundell, Rawlings \\& Willott (1999) showed that any radio source evolution involving a decrease in luminosity with time would, at the highest redshifts, result in a bias towards young sources in flux density limited samples. Since this effect is only important at $z>2-3$, it is not relevant to the analysis presented in this paper.  Several GPS and CSOs (eg. 0108+388; Baum et al. 1990) exhibit low level, steep spectrum, extended emission on arcsecond scales, which seem to be relics of much older radio-activity. These objects are often classified as being intermittent or re-occurent, and therefore not as young objects. However, the components related to their gigahertz-peaked spectra and CSO morphologies are certainly young. The presence of faint relic emission only indicates that the active nucleus has been active before, and may constrain the typical timescale and frequency of such events. Based on the current knowledge of the formation of massive black-holes in the centers of galaxies, it is unlikely that the central engine itself is young, but just the radio source (Richstone et al. 1998).  It is unclear whether all young sources actually evolve to large extended objects. Some, or even the majority, may be short-lived phenomena due to a lack of significant fuel (Readhead 1994). The possible existence of these objects can have a large influence on the source statistics of young radio sources. ",
        "conclusions": ""
    },
    "0002/astro-ph0002289_arXiv.txt": {
        "abstract": "Images of Hubble 12 were obtained with the HST/NICMOS instrument in the F110W, F164N, and F166N filters of NIC1, and the F160W, F187N, F190N, F212N, F215N filters of NIC2. The images show the structure of the inner ``torus\" and lobes much clearer than previous ground-based images.  In particular, the [Fe~II] lobes are clearly resolved and shown to be distinct from the H$_2$ emission structures. Apparent changes in the inner ionized and neutral bipolar shell implies periodic mass loss or changes in stellar wind shape or direction. The H$_2$ in the ``eye\" is radiatively excited and shows a complex morphology that suggests that several mass ejection events are responsible for producing this structure. The position angle of the  H$_2$ and [Fe~II] lobes differ, indicating a possible precession of the ejection axis. ",
        "introduction": "Images of Hubble 12 (Hb 12) were obtained with the HST/NICMOS instrument on 13 Nov 1997.  The images in the F110W, F164N, and F166N were taken with NIC1, and the F160W, F187N, F190N, F212N, F215N images were obtained with NIC2. The MULTIACCUM mode and spiral dither patterns were used.  The  NICMOS CALNICA and CALNICB pipelines were used to reduce the data.  In general a few dither sets were obtained in each filter; these were aligned and averaged to produce the final images. The bright central star makes the region near the core difficult to see clearly, and adds a number of image artifacts.  The lines that run approximately from corner to corner are part of the diffraction pattern of the instrument. The horizontal and vertical lines that run through the star are array artifacts caused when the central star saturates in the long integrations.  In future reductions we will attempt to subtract the point source in the core to better study the region around it.  \\begin{figure} \\plottwo{hora_fig1a.ps}{hora_fig1b.ps} \\caption{Hubble 12, in the F110W (left) and F164N filters (right).  The field size for each is $\\sim$ 13 arcsec square.  The orientation of the images shown in the F110W image is the same for all images presented here.  The F110W filter includes contributions from the bright Paschen $\\beta$ line and line emission from H$_2$, [Fe~II], and He~I lines within the bandpass. The F164N filter samples the [Fe~II] line at 1.64 $\\mu$m.  } \\end{figure}  \\begin{figure} \\plottwo{hora_fig2a.ps}{hora_fig2b.ps} \\caption{Hubble 12, in the F160W (left) and F187N filters (right).  The field size shown in each image is $\\sim$ 19 arcsec square. The bandpass of the F160W filter includes lines from the Brackett series of H~I and line emission from H$_2$, [Fe~II], and He~I lines as well.  The narrowband F187N filter samples the Paschen $\\alpha$ (H~I) line. } \\end{figure} ",
        "conclusions": "Hb 12 has been notable primarily because it represents one of the clearest cases known of UV excited near-IR fluorescent H$_2$ emission (Dinerstein et al.\\ 1988; Ramsay et al.\\ 1993, Hora \\& Latter 1996, Luhman \\& Rieke 1996). Dinerstein et al.\\ had mapped the inner structure and found it to be elliptical surrounding the central star;  the deep H$_2$ images in Hora \\& Latter (1996)  showed the faint bipolar lobes extending N-S, and the torus or ``eye\"-shaped strucure at the base of the lobes around the central star. The H$_2$ line ratios observed in the torus were in excellent agreement with predictions by theoretical H$_2$ fluorescence calculations (see also Luhman \\& Rieke 1996). Hora \\& Latter also detected [Fe~II] line emission at 1.64 \\micron\\ in a position along the edge of the shell, but not at the H$_2$ line peak emission location to the E of the central star.  \\begin{figure} \\plottwo{hora_fig3a.ps}{hora_fig3b.ps} \\caption{Hubble 12, in the F212W (left) and F215N filters (right).  The field size shown in each image is approximately 19 arcsec square. The F212N filter samples the H$_2$ line at 2.12 $\\mu$m, the F215N filter measures the nearby continuum (as well as a small contribution from the Brackett $\\gamma$ feature at 2.16 $\\mu$m.} \\end{figure}  Previous HST-WPC2 imaging by Sahai \\& Trauger (1998) in H$\\alpha$ showed the inner structure to have an ``hourglass\" shape, and a small bipolar structure in the core region, with lobes roughly E-W within a few tenths of an arcsec from the star. Welch et al. (1999a,b) obtained ground-based images in the [Fe~II] line and nearby continuum and found that the line emission was also distributed along the inner hourglass nebula. The HST images presented here show the symmetry axes of the hourglass and the H$_2$ eye and bipolar nebula differ in their alignment by $\\sim$ 5$^\\circ$. A comparison of these images with the inner bipolar structure found by Sahai \\& Trauger shows that its alignment differs by $\\sim 20^\\circ$ from the hourglass and outer H$_2$ lobes.  The different orientation of the structures suggests that the central source may be precessing between discrete outflow events.  Also, the structures seen in the H$_2$ image indicate other possible outflow events and remnants of other bipolar hourglass nebulae.  Hb~12 may therefore be another example of a PN with multiple nested bipolar bubbles.  The inner hourglass is bright in the Paschen $\\alpha$ and [Fe~II] lines, but the outline of the ``eye'' appears only in the H$_2$ and the wide bandpass filters (in continuum plus H$_2$ line emission).  This implies that the regions where only H$_2$ is detected are somehow shielded from what has ionized the inner hourglass. There is no evidence  for [Fe~II] emission from other regions in the nebula, and the density of the inner hourglass does not seem sufficient to provide effective shielding of the other regions from radiation from the central star, which might suggest shock excitation in an interacting wind.  This hypothesis must be confirmed with investigation of the velocity structure in the nebula. The strong influence of FUV photons elsewhere in this object argues that the [Fe~II] emission is excited by FUV photons in the PDR. We will be investigating this possibility through detailed chemical modeling.  We obtained high-resolution IR spectra of Hb~12 in the H$_2$ and [Fe~II] lines using CSHELL at the IRTF (Kelly, Hora, \\& Latter 1999) which indicates that the N lobe is inclined towards us.  We also have additional medium-resolution spectra of the faint extended H$_2$ lobes to determine the excitation properties of the outer nebula.  We are in the process of analyzing these data along with the optical and IR imagery to understand the structure of this interesting and complex nebula."
    },
    "0002/astro-ph0002076_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "\\label{conc.sec} Experimental studies of EAS characteristic like the depth of the shower maximum $X_{max}$, the elongation rate $dX_{max}/dlog_{10}E$ and the attenuation length $\\Lambda$ are of particular importance, since they map rather directly basic features of the hadronic interaction. Strictly, however, the interpretation of these quantities in terms of hadronic cross sections cannot bypass the necessity of detailed calculations of the shower development. Nevertheless these type of EAS quantities, if compared with Monte Carlo simulation results, provide stringent tests of the interaction model ingredients of the simulations. \\\\ The recent results of various experimental installations are sufficiently accurate to enable relevant studies of this kind, and combining the data from arrays situated on different altitudes (like MAKET ANI and KASCADE) allows a large span in the atmospheric slant depth for reconstructing the development of the charge particle size. In fact such studies, if using a sufficiently large data sample, could  be continued in a more detailed manner by separating the muon component and taking into account the deviations from the exponential shape of the cascade decline. The penetrating muon component contributes with smaller attenuation to the development of the considered charged particle component, but hardly with an exponential degrading (according to eq.\\ref{cascade}). Actually by use of methods in progress to isolate different primary groups (\"pure nuclear beams\") of the size spectra \\cite{beam2,gagik99}, these kind of interaction studies would get of extreme interest.  \\vspace*{3.5ex} {\\noindent \\bf \\Large Acknowledgment \\\\} \\vspace*{-.5ex}  {\\it \\noindent This publication is based on experimental results of the ANI collaboration. The MAKET ANI detector installation has been set up as collaborative project of the Yerevan Physics Institute, Armenia and the Lebedev Institute, Moscow. The continuous contributions and assistance of the Russian colleagues in operating the detector installation and in the data analyses are gratefully acknowledged. In particular, we thank S. Nikolski and V. Romakhin for their encouraging interest and useful discussions. \\\\ First perspectives of combined considerations of the KASCADE and ANI experimental data have been discussed in 1998 during the ANI-98 workshop in the cosmic ray observatory station Nor-Amberd of Mt.Aragats (Armenia). The MAKET ANI group would like to thank the German colleagues for stimulating discussions and encouragement, in particular H. Rebel for his numerous valuable comments and interesting suggestions to the topic of this paper. We acknowledge the useful discussions with K.-H. Kampert, H. Klages and R. Glasstetter. The suggestions of S. Ostapchenko are highly appreciated. \\\\ The work has been partly supported by the research grant No.96-752 of  the Armenian Government and by the ISTC project A116. The assistance of the Maintenance Staff of the Aragats Cosmic Ray Observatory in operating the MAKET ANI installation is highly appreciated.}  \\vspace*{-0.3cm}"
    },
    "0002/hep-ph0002147.txt": {
        "abstract": "A global analysis of all the available atmospheric neutrino data is presented in terms of neutrino oscillations in the $\\nu_\\mu \\to \\nu_\\tau$ and $\\nu_\\mu \\to \\nu_s$ channels, where \\ns denotes a sterile neutrino.  We perform our analysis of the contained events data as well as the upward--going neutrino--induced muon fluxes. In addition to the previous data samples of Frejus, Nusex, IMB and Kamioka experiments, we include the full data set of the 52 kton-yr of Super--Kamiokande, the recent 4.6 kton-yr contained events of Soudan2 and the results on upgoing muons from the MACRO and Baksan detectors.  From the statistical analysis it emerges that the $\\nu_\\mu \\to \\nu_\\tau$ channel provides the best agreement with the combined data, with a best fit point of $\\sin^2(2\\theta) = 0.99$ and $\\Delta m^2 = 3.0 \\times 10^{-3}$ eV$^2$. Although somehow disfavoured, the $\\nu_\\mu \\to \\nu_s$ channels cannot be ruled out on the basis of the global fit to the full set of observables. ",
        "introduction": "Together with the solar neutrino problem \\cite{Bahcall:1999rt} the atmospheric neutrino anomaly constitutes the second evidence for physics beyond the Standard Model. Indeed a large number of observations \\cite{Frejus,Nusex,IMB,Soudan,Soudan2,Kamiokande,skcont,skup,sk99,MACRO,Baksan} have been performed concerning electron-- and muon--neutrino fluxes produced by hadronic showers initiated by cosmic--ray interactions in the upper atmosphere. Apart from the first iron--calorimeter detectors \\cite{Frejus,Nusex}, all experiments, which also entail different detection techniques, have steadily reported a deficit of the collected number of events with respect to the theoretical expectations \\cite{fluxes1,fluxes2}.  Although the knowledge of the fluxes of atmospheric neutrinos is affected by uncertainties which range from about 20\\% to 30\\%, the expected ratio $R(\\mu/e)$ of the muon neutrino ($\\nu_\\mu + \\bar{\\nu}_\\mu$) over the electron neutrino flux ($\\nu_e+\\bar{\\nu}_e$) is known with much better confidence, since the uncertainties associated with the absolute fluxes largely cancel out. % It is fair to ascribe an overall uncertainty less than about 5\\% to the calculated $R(\\mu/e)$ ratio. Since this ratio is measured to be substantially smaller than the expectations, one faces an anomaly which only seems possible to account for in terms of non--standard neutrino properties.  Super--Kamiokande high statistics observations~\\cite{skcont,skup} indicate that the deficit in the total ratio $R(\\mu/e)$ is due to the neutrinos arriving in the detector at large zenith angles. They also show that the $e$-like events do not present any compelling evidence of a zenith-angle dependent suppression, while the $\\mu$-like event rates are substantially suppressed at large angles. Different explanations to these features have been proposed and discussed in the literature \\cite{others1,others2}.  The simplest and most direct possibility is represented by the oscillation of muon neutrinos $\\nu_\\mu$ into either a $\\nu_\\tau$ or a sterile neutrino $\\nu_s$ \\cite{atm98,atmo98,skos}.  The oscillation hypothesis provides a very good explanation for this smaller-than-expected ratio, which is also simple and well-motivated theoretically.  In this paper we present our updated analysis of all the available data on atmospheric neutrinos in terms of neutrino oscillation. We include in the analysis all the existing experimental results obtained so far. In addition to the previous data samples of Frejus\\cite{Frejus}, Nusex\\cite{Nusex}, IMB\\cite{IMB} and Kamioka\\cite{Kamiokande} experiments, we also consider the recent Soudan2\\cite{Soudan2} data which refers to an exposure of 4.6 kton-yr, the full data set of the 52 kton-yr of Super--Kamiokande\\cite{sk99}, including both contained events and upgoing muon data, and finally the results on upgoing muons reported from the MACRO\\cite{MACRO} and Baksan\\cite{Baksan} experiments. We critically discuss the analysis of the various individual data sets. Moreover we consider the combined information that can be derived from all the experimental evidences so--far obtained. We hope this will contribute to a more complete understanding of the atmospheric neutrino anomaly in the framework of a global analysis based on a self--consistent theoretical calculation of the event rates. This has the advantage of calibrating all relevant elements in the analysis, such as theoretical atmospheric neutrino fluxes, in a consistent way. This allows different experiments to be compared in a meaningful way. A new element which we now add to our previous investigations\\cite{atm98} is the inclusion of the upgoing muon data and the corresponding theoretical flux determinations\\cite{sk99,MACRO,Baksan}. We anticipate that from the statistical analysis the $\\nu_\\mu \\to \\nu_\\tau$ emerges as the oscillation channel which provides the best agreement with the combined data. The $\\nu_\\mu \\to \\nu_s$ channels, although slightly disfavoured, cannot be statistically ruled out on the basis of the global fit to the full set of observables.  The outline of the paper is the following: in Sect. II we briefly recall the theoretical calculation of the event rates for contained events and upgoing muon fluxes, as well as the calculation of the oscillation probabilities. Sect. III discusses the statistical approach of our analysis and reports on the results of the fits. Sect. IV presents a discussion of the results and the conclusions. We include more details on the statistical analysis in an Appendix. ",
        "conclusions": "In this paper we have performed a global analysis of the atmospheric neutrino data in terms of neutrino oscillations. We have compared the relative statistical relevance of the active-active and active-sterile channels as potential explanations of the atmospheric neutrino anomaly.  In the analysis we have included, for the contained events, the latest data from Super--Kamiokande, corresponding to 52 kton-yr, together with all other available atmospheric neutrino data in the sub-GeV and multi-GeV range. Specifically, we included the data sets of the Frejus, Nusex, IMB, Soudan2 and Kamiokande experiments.  Our analysis also includes the results on neutrino induced upgoing muons from the Super--Kamiokande 52 kton-yr sample, from MACRO and from Baksan experiments.  We have determined the best-fit neutrino oscillation parameters and the resulting allowed regions in $\\sin^2(2\\theta)$ and $\\Delta m^2$ for $\\nu_\\mu \\to \\nu_X$ ($X=\\tau,s$) channels. For oscillations into sterile neutrinos we have considered both positive and negative $\\Delta m^2$, since the two cases differ in the matter effect for neutrinos propagation in the Earth.  The results of the combined analysis indicate that the channel $\\nu_\\mu \\to \\nu_\\tau$ gives a better fit to the data than oscillations into sterile neutrinos: $\\chi^2_{min} = $ 74/(73 d.o.f.) (45 \\% CL) for the active case versus $\\chi^2_{min} = $ 90/(73 d.o.f.) (8.5 \\% CL) for oscillations into sterile neutrinos with $\\Delta m^2 <0$ and $\\chi^2_{min} = $ 86/(73 d.o.f.) (14 \\% CL) for $\\nu_\\mu \\to \\nu_s$ with $\\Delta m^2 >0$. Since for some experiments we obtain large $\\chi^2_{min}$ values also for neutrino oscillations, values not clearly better than the analysis for the standard model case, we decided to perform additional analyses by removing these ``negative-result experiments''.  When excluding Frejus, Nusex and Baksan data points, we do not obtain an improvement for the sterile cases, while for the active case the CL is increased to 58 \\% ($\\chi^2_{min} = $ 58/(61 d.o.f.)). When also MACRO is removed, we obtain higher CL also for the sterile cases, but the $\\nu_\\mu \\to \\nu_\\tau$ hypothesis remains as the best option: $\\chi^2_{min} = $ 41/(51 d.o.f.) (84 \\% CL) for the active case, $\\chi^2_{min} = $ 51/(51 d.o.f.) (47 \\% CL) for $\\nu_\\mu \\to \\nu_s$ with $\\Delta m^2 <0$ and $\\chi^2_{min} = $ 50/(51 d.o.f.)  (51 \\% CL) for $\\nu_\\mu \\to \\nu_s$ with $\\Delta m^2 >0$. We thus conclude that the quality of the global description of the atmospheric neutrino data in terms of neutrino oscillation is better for the $\\nu_\\mu \\to \\nu_\\tau$ channel, although oscillations into $\\nu_\\mu \\to \\nu_s$ cannot be statistically ruled out, on the basis of the global fit to the full set of observables.  We have also presented a sample of predicted zenith-angle distributions for the best-fit points corresponding to the various oscillation channels.  Specifically, we showed the angular distribution corresponding to the four Super--Kamiokande data sets (sub-GeV, multi-GeV, stopping muons and through-going muons) and the angular distributions for upgoing muons at MACRO and Baksan.  By using the zenith-angle distribution expected for contained events and up-going muon data at Super-Kamiokande in the presence of oscillations, we have compared the relative goodness of the three possible oscillation channels. This allowed one to understand why the $\\nu_\\mu \\to \\nu_\\tau$ channel gives a better description than the $\\nu_\\mu \\to \\nu_s$ cases. On the one hand, due to matter effects, one finds that for the sterile case the up--down asymmetry in the multi-GeV sample is slightly smaller than observed.  Moreover, also due to matter effects, the upgoing-muon distributions in the case of $\\nu_\\mu \\to \\nu_s$ are flatter than for $\\nu_\\mu \\to \\nu_\\tau$, while the data show a slightly steeper angular dependence which can be better described by $\\nu_\\mu \\to \\nu_\\tau$.  To conclude, we compare our determinations of the allowed neutrino oscillation parameters from the analyses of the atmospheric neutrino data with the expected sensitivities of future long-baseline experiments such as K2K\\cite{k2k} and MINOS\\cite{minos} (Fig.~\\ref{regionglobal}).  We notice that, for the oscillations $\\nu_\\mu \\to \\nu_\\tau$, K2K can cover most of the 90\\% CL allowed region while the MINOS test of NC/CC is sensitive to the complete 99\\% CL region of oscillation parameters. For oscillations into sterile neutrinos, a situation where only a CC disappearance test can be performed at long baseline experiments, K2K can cover a portion of the region allowed at 90\\% CL, while the full 99\\% CL allowed regions are completely accessible to the MINOS experiment."
    },
    "0002/astro-ph0002526_arXiv.txt": {
        "abstract": "We have obtained moderate resolution ($R$ = a few thousand) spectra of the $NaI\\lambda\\lambda$5890,5896 ($NaD$) absorption-line in a sample of 32 far-IR-bright starburst galaxies. In 18 cases, the $NaD$ line in the nucleus is produced primarily by interstellar gas, while cool stars contribute significantly in the others. In 12 of the 18 ``interstellar-dominated'' cases the $NaD$ line is blueshifted by over 100 km s$^{-1}$ relative to the galaxy systemic velocity (the ``outflow sources''), while no case shows a net {\\it redshift} of more than 100 km s$^{-1}$. The absorption-line profiles in these outflow sources span the range from near the galaxy systemic velocity to a maximum blueshift of $\\sim$ 400 to 600 km s$^{-1}$. The outflow sources are galaxies systematically viewed more nearly face-on than the others. We therefore argue that the absorbing material consists of ambient interstellar material that has been entrained and accelerated along the minor axis of the galaxy by a hot starburst-driven superwind. The $NaD$ lines are optically-thick, but indirect arguments imply total Hydrogen column densities of $N_H \\sim$ few $\\times 10^{21}$ cm$^{-2}$. This implies that the superwind is expelling matter at a rate comparable to the star-formation rate. This outflowing material is evidently very dusty: we find a strong correlation between the depth of the $NaD$ profile and the line-of-sight reddening. Typical implied values are $E(B-V)$ = 0.3 to 1 over regions several-to-ten kpc in size. We briefly consider some of the potential implications of these observations. The estimated terminal velocities of superwinds inferred from the present data and extant X-ray data are typically 400 to 800 km s$^{-1}$, are independent of the galaxy rotation speed, and are comparable to (substantially exceed) the escape velocities for $L_*$ (dwarf) galaxies. The resulting selective loss of metals from shallower potential wells can establish the mass-metallicity relation in spheroids, produce the observed metallicity in the intra-cluster medium, and enrich a general IGM to of-order 10$^{-1}$ solar metallicity. If the outflowing dust grains can survive their journey into the IGM, their effect on observations of cosmologically-distant objects would be significant. ",
        "introduction": "By now, it is well-established that galactic-scale outflows of gas (sometimes called `superwinds') are a ubiquitous phenomenon in the most actively star- forming galaxies in the local universe (Heckman, Lehnert, \\& Armus 1993; Dahlem 1997; Bland-Hawthorn 1995). They are  powered by the energy deposited in the interstellar medium by massive stars via supernovae and stellar winds. Over the history of the universe, outflows like these may have polluted the intergalactic medium with metals (e.g. Giroux \\& Shull 1997) and dust (Alton, Davis, \\& Bianchi 1999; Aguirre 1999a,b), heated and polluted the intracluster medium (e.g. Gibson, Loewenstein, \\& Mushotzky 1997; Ponman, Cannon, \\& Navarro 1999), and may have established the mass-metallicity relation and radial metallicity gradients in galactic spheroids  (e.g. Carollo \\& Danziger 1994). However, the astrophysical relevance of superwinds can not be reliably assessed without first understanding their  physical, dynamical, and chemical properties. To date, most of the pertinent information has come from observations of  the X-ray emission produced by the hot gas (e.g. Dahlem, Weaver, \\& Heckman 1998) or the optical line-emission produced by the warm gas (e.g. Lehnert \\& Heckman 1996a).  In the present paper, we take a complementary approach, and discuss an extensive body of new data that probes  the outflowing gas via its interstellar absorption-lines. This technique has some important advantages. First, since the gas is seen in absorption against the background starlight, there is no possible ambiguity as to the sign (inwards or outwards) of any radial flow that is detected. Second, the strength of the absorption will be related  to the column density of the gas. In contrast, the X-ray or optical surface-brightness of the emitting gas is proportional to the emission-measure. Thus, the absorption-lines more fully probe the whole range of gas densities in the outflow, rather than being strongly weighted in favor of the densest material (which may contain relatively little mass).  Finally, provided that suitably- bright background sources can be found, interstellar absorption-lines can be used to study outflows in high-redshift galaxies where the associated X-ray or optical {\\it emission} may be undetectably faint. This promise has already been realized in the case of the 'Lyman Dropout' galaxies, where the kinematic signature of outflows is clear in their rest-frame UV spectra (Franx et al 1997; Pettini et al 1998, 1999).  A few pioneering studies have already detected interstellar absorption-lines from superwinds in local starburst galaxies. Phillips (1993) discussed spatially-resolved optical spectroscopy of the $NaI$ ``D'' doublet in NGC 1808, showing that an outflow of gas at velocities of  up to 700 km s$^{-1}$ could be traced over a region several kpc in size, coincident with a region of extra- planar dust plumes. Several recent papers (Lequeux et al. 1995; Heckman \\& Leitherer 1997; Sahu \\& Blades 1997; Kunth et al. 1998; Gonzalez-Delgado et al 1998a) have detected blueshifted interstellar absorption-lines in $HST$ and $HUT$ UV spectra of a handful of starburst galaxies, implying outflows of metal- bearing gas at velocities of 10$^2$ to 10$^3$ km s$^{-1}$.  Most of the strong resonance lines of cosmically-abundant ions are found in the UV (e.g. Morton 1991; Savage \\& Sembach 1996), and so must be studied with $HST$ or $FUSE$ in local starbursts. In the present program, we have instead exploited the relatively greater sensitivity and availability of ground-based telescopes at visible wavelengths to study a large sample of starbursts/superwinds using the $NaI$ doublet at $\\lambda$$\\lambda$5890,5896 \\AA. In a few cases, we have also observed the $KI$ $\\lambda$$\\lambda$7665,7699 \\AA\\ doublet, since it probes gas with nearly the same ionization state as $NaI$, but is more likely to be optically thin (the $K$ abundance is down from $Na$ by a factor of 15, while the two doublets have similar oscillator strengths). The ionization potentials of $NaI$ and $KI$ are only 5.1 eV and 4.3 eV respectively, so these species should primarily probe the $HI$ and $H_2$ ISM phases. Observations in vacuum-UV will be required to study the hotter and more highly ionized gas in absorption. ",
        "conclusions": "We have discussed the results of moderate-resolution ($R$ = a few thousand) spectroscopy of the $NaI\\lambda\\lambda$5890,5896 ($NaD$) absorption-line in a sample of 32 far-IR-selected starburst galaxies. These galaxies were selected from either the far-IR-warm sample of Armus, Heckman, \\& Miley (1989) or the edge-on sample of Lehnert \\& Heckman (1995), and together span a range from 10$^{10}$ to few $\\times$ 10$^{12}$ L$_{\\odot}$ in IR luminosity. We found that the stellar contribution to the $NaD$ absorption-line is negligible ($<$10\\%) in some objects, but significant ($\\sim$ 70\\%) in others. We have thus divided our sample into 18 interstellar-dominated (``ISD'') objects ($<$ 30\\% stellar contribution) and 14 strong-stellar-contamination (``SSC'') objects ($>$ 40\\% stellar contribution).  The $NaD$ line lies within 70 km s$^{-1}$ of $v_{sys}$ in all the SSC objects (consistent with a predominantly stellar origin). The $NaD$ lines in the SSC nuclei are about 0.2 dex narrower than expected for dynamics of the old stellar population in the bulges of normal galaxies of similar disk rotation speed and Hubble type. Thus, dynamically ``cold'' material (red supergiants and/or interstellar gas) in the inner part of the starburst makes a significant contribution to the observed $NaD$ line in these nuclei.  The kinematics of $NaD$ line are markedly different in the ISD objects. The $NaD$ line is blueshifted by $\\Delta$$v >$ 100 km s$^{-1}$ relative to the galaxy systemic velocity in 12 of the 18 cases (the ``outflow sources''), and the outflow can be mapped over a region of a few-to-ten kpc in size. In contrast, no objects in our sample showed a net {\\it redshift} in $NaD$ of more than 100 km s$^{-1}$. The outflow sources are galaxies systematically viewed more nearly face-on than the other galaxies in our sample: 69\\% of the galaxies with a ratio of semi-major to semi-minor axes $a/b \\leq$ 2.0 show $NaD$ outflows, while this is true for only 6\\% of the flatter (more highly inclined) galaxies. This is consistent with the absorbing material being accelerated out along the galaxy minor axis by a bipolar superwind. The absorbing material typically spans the velocity range from near the galaxy systemic velocity ($v_{sys}$) to a maximum blueshift of 300 to 700 km s$^{-1}$. We therefore suggest that the outflowing superwind ablates the absorbing gas from ambient clouds at $\\sim v_{sys}$, and then accelerates it up to a terminal velocity similar to the wind speed. We found no correlation between the widths of the H$\\alpha$ emission-line and the $NaD$ absorption-line subsamples. Evidently, the dynamics of the more tenuous absorbing gas is largely decoupled from that of the dense (high emission-measure) gas that provides most of the nuclear line-emission.  The ratio of the equivalent widths of the two members of the $NaD$ doublet ($R$) ranges from 1.1 to 1.7 in the ISD sample, implying that the doublet is optically-thick. However, $R$ does not correlate with the residual relative intensity at the ``bottom'' of the stronger $\\lambda$5890 line profile ($I_{5890}$), which ranges from 0.14 (nearly black) to 0.7. Thus, the optically-thick gas does not fully cover the emitting stars (covering factor $\\sim$ 1 - $I_{5890}$). The observed equivalent width of the $NaD$ line is then set by the product of velocity dispersion and covering factor for the absorbing gas, and we showed that the latter quantity is the dominant one. Using two variants of the classic doublet-ratio technique, we estimated that the $NaI$ column densities are $logN_{NaI}$ = 13.5 to 14 cm$^{-2}$. This is roughly consistent with column densities measured in a few cases for $KI$ using the optically-thin $\\lambda$$\\lambda$7665,7699 \\AA\\ doublet (assuming a solar $NaI/KI$ ratio). The total gas columns are uncertain, but the empirical correlation between $N_{NaI}$ and $N_H$ in the ISM of the Milky Way implies $N_H \\sim$ few $\\times$ 10$^{21}$ cm$^{-2}$.  We found a strong correlation in the ISD sample between the reddening of the observed stellar continuum and the depth of the $NaD$ absorption-line, and a significant but weaker correlation of the line-depth with the reddening of the Balmer emission-lines. Evidently, the gas responsible for the $NaD$ absorption is very dusty. The typical implied reddening is $E(B-V) \\sim$ 0.3 to 1 magnitudes over regions several-to-ten kpc in size. For a normal dust-to-gas ratio, the corresponding column densities are $N_H \\sim$ few $\\times$ 10$^{21}$ cm$^{-2}$ (in agreement with the above estimate).  The inferred column densities and measured outflow velocities and sizes imply that the typical mass and kinetic energy associated with the absorbing gas is of-order 10$^9$ M$_{\\odot}$ and 10$^{56}$ erg, respectively. The estimated outflow rates of mass and energy are typically 10 to 100 M$_{\\odot}$ per year and 10$^{41}$ to 10$^{42}$ erg s$^{-1}$. The mass outflow rates are comparable to the estimated star-formation rate, and much larger than the rate at which massive stars are returning mass to the ISM. Thus, powerful starbursts can eject as much gas as is being converted into stars, and most of this gas is ambient material that has been ``mass-loaded'' into the hot gas returned directly by supernovae and stellar winds. The energy outflow rates in the absorption-line gas are of-order 10$^{-1}$ of the rate at which massive stars supply mechanical energy. Most of the energy returned by massive stars probably resides in the kinetic and thermal energy of the much hotter X-ray-emitting gas. We showed that the overall properties of the absorbing gas in the outflow sources can be easily reproduced in the context of simple analytic estimates for the properties of interstellar clouds accelerated by the ram pressure of the hot high-speed wind seen via its X-ray emission. Detailed hydrodynamical simulations of galactic winds, while still missing some essential physics, also predict the observed properties of the cool absorbing gas.  We have discussed the implications of our results for the chemical evolution of galaxies and the intergalactic medium. The estimates derived for $v_{term}$ using the $NaD$ line in the outflow sources agree reasonably well with the outflow speeds implied for an adiabatic wind ``fed'' by hot gas whose temperature is measured by the observed X-ray-emitting gas. The typical implied values are 300 to 800 km s$^{-1}$, and are independent of the rotation speed of the ``host galaxy'' over the range $v_{rot}$ = 30 to 300 km s$^{-1}$, confirming and extending the result in Martin (1999) based on X-ray data alone. This strongly suggests that the outflows selectively escape the potential wells of the less massive galaxies. We considered a simple model based on Lynden-Bell (1992) in which the fraction of starburst-produced metals that are retained by a galaxy experiencing an outflow is proportional to the galaxy potential-well depth for galaxies with $v_{esc} < v_{term}$, and asymtotes to full retention for the most massive galaxies ($v_{esc} > v_{term}$). For $v_{term}$ in the range we measure, such a simple prescription can reproduce the observed mass-metallicity relation for elliptical galaxies and deposit the required amount of observed metals in the intra-cluster medium. If the ratio of ejected metals to stellar spheroid mass is the same globally as in clusters of galaxies, we predicted that the present-day mass-weighted  metallicity of an intergalactic medium with $\\Omega_{igm}$ = 0.015 will be $\\sim$ 1/6 solar (see also Renzini 1997).  We have summarized the evidence that starbursts are ejecting significant quantities of dust, emphasizing the results from the present paper. {\\it If} this dust can survive a trip into the intergalactic medium and remain intact for a Hubble time, we estimated that the upper bound on the global amount of intergalactic dust is $\\Omega_{dust}$ $\\sim$ 10$^{-4}$. While this is clearly an upper limit, it is a cosmologically interesting one: Aguirre (1999a,b) argues that dust this abundant could in principle obviate the need for a positive cosmological constant, based on the Type Ia supernova Hubble diagram.  Finally, given the mounting evidence for a connection between starbursts and the Seyfert phenomenon, we have suggested that outflows like those studied here may account for some (but not all) aspects of the ``associated absorption'' seen in type 1 Seyfert nuclei."
    },
    "0002/astro-ph0002183_arXiv.txt": {
        "abstract": "This paper presents the first results of a survey for bright quasars ($V < 14.5$ and $R<15.4$) covering the North Hemisphere at galactic latitudes $|b|>30$. The photometric database is derived from the GSC and USNO catalogs. Quasars are identified on the basis of their X-ray emission measured in the ROSAT All Sky Survey. The surface density of quasars brighter than 15.5 mag turns out to be $10 \\pm 2 \\cdot 10^{-3} deg^{-2}$, about 3 times higher than that estimated by the PG survey. The quasar optical Luminosity Function (LF) at $0.04 < z \\le 0.3$ is computed and shown to be consistent with a Luminosity Dependent Luminosity Evolution of the type derived by La Franca and Cristiani (1997) in the range $0.3 < z \\le 2.2$. The predictions of semi-analytical models of hierarchical structure formation agree remarkably well with the present observations. ",
        "introduction": "Quasar surveys provide basic information for the understanding of a number of astrophysical, cosmological and cosmogonical issues: the formation and evolution of galactic structures, the physics of the AGN phenomenon, the UV and X-ray backgrounds.  The general behaviour of the QSO optical luminosity function (OLF) is well established in the redshift interval $0.3<z<2.2$ for which color techniques provide reliable selection criteria (see \\cite{HS:90,boyle92,HF94}, for a review of the subject). A pure luminosity evolution (PLE) appears to reasonably describe tha data in the interval $0.6<z<2.2$. In the range $0.3<z<0.6$ the OLF appears to be flatter than observed at higher redshifts (\\cite{pippa92}), requiring a luminosity dependent luminosity evolution (LDLE, \\cite{lf97} LC97). This departure from a PLE provides an interesting clue for the physical interpretation of the QSO evolving population (\\cite{caval97}, LC97). The present observational evidence, however, relies on a relatively small number of objects: in the $0.3<z<0.6$ range for $M_B<-25$, $32$ QSOs are observed by LC97 instead of the $19$ expected from the best-fit PLE model.  Analogous results by K\\\"ohler et al. (1997) and Goldschmidt and Miller (1998) are likewise based on very small samples.  To provide a statistically solid basis for the LDLE pattern and to investigate whether such a trend persists (and possibly becomes more evident) at redshifts lower than 0.3, we decided to carry out a new large-area survey of quasars at bright optical fluxes. A typical apparent magnitude for a $M_B = -24$ QSO at $z\\sim 0.1$ is in fact $B \\sim 14.5$ and the surface density of these objects, according to previous surveys (e.g. the BQS, \\cite{SeG83}), is expected to be less than $10^{-3}$ per sq.deg., requiring an efficient selection criterion and the coverage of a significant fraction of the whole sky in order to collect a meaningful sample.  In Sect.2 we describe the photometric database from which optical fluxes have been derived, in Sect.3 the criteria followed to select the candidates, in Sect.4 the spectroscopic follow-up, in Sect.5 the derived quasar counts and the optical luminosity function and in Sect.6 a few consequences for model scenarios. ",
        "conclusions": "Franceschini et al. (1994) and La Franca et al. (1995) have shown that the QSO emission in the soft-X and at visible wavelengths scales linearly.  A confirmation of the constant ratio $L_X/L_O$ and independence both from $L_O$ and from $z$ comes from Yuan et al. (1998). Boyle et al. (1994) derived a XLF comparable to the QSO OLF known at that epoch with an evolutionary rate $L_X(z)=L_X(0)(1+z)^{3.25}$, similar to the optical one. These works favor a scenario in which essentially the same QSO population is observed, both in the soft-X and in the optical. In this way the flattening of the OLF observed in the present survey should be reflected in a flattening in the corresponding soft-X LF. Indeed Miyaji et al. (1999) show that the bright part of the 0.5-2 keV LF flattens with decreasing redshift from a value $1+\\beta \\simeq 3.5$ at $0.4 < z < 0.8$ to $1+\\beta \\simeq 2.6$ at $0.015 < z < 0.2$, which is similar to what we observe in the optical: $1+\\beta \\simeq 3.7$ at $0.6 < z < 1.0$ and $1+\\beta \\simeq 3.1$ at $0.04 < z \\le 0.3$.  The decline of the space density of quasars from a peak around redshift 2 to the present epoch has been modeled by several authors in the framework of hierarchical theories of structure formation (\\cite{cattaneo:99,HM:99,KH:99,MSD:99}). In particular Kauffmann and Haehnelt (1999, KH99) have attempted to reproduce quantitatively the evolution of the quasar number density incorporating a scheme for the growth of massive black holes (MBHs) into semi-analytical models following the evolution of galaxies in CDM-dominated scenarios. Together with the decrease in the merging rates and in the amount of gas available to fuel the MBHs, which are built-in features of the semi-analytic models, KH99 assume an increase of the timescale for gas accretion in order to reproduce the steep decline in the number density of quasars from $z \\sim 2$ to $z=0$. Other authors have followed similar recipes assuming a decreasing mass accretion (\\cite{HM:99}), or a decreasing efficiency of the accretion (\\cite{cattaneo:99}), or a delayed quasar activity with respect to the dynamical formation of the halos with a longer delay for smaller halos (\\cite{MSD:99}).  As can be seen from Fig.~17 of KH99, the semi-analytical models have difficulties in reproducing the steep decrease of the QSO density at low redshift that is commonly measured (\\cite{HS:90}). The most promising scenario is a $\\Lambda$CDM, in which the accretion timescale, $t_{\\rm acc}$, is assumed to vary in the same way as the host galaxy dynamical time ($t_{\\rm acc} \\propto [0.7 + 0.3(1+z)^3]^{-1/2}$). This model is able to reproduce the evolution of the galaxy LF and of the cold gas content of galaxies, but is apparently predicting a too slow quasar decline. The present data significantly reduce this disagreement in the sense that the higher quasar space density measured in our survey corresponds fairly well to the $\\Lambda$CDM semi-analytical predictions, as shown in Fig.~\\ref{semi_anal}. \\placefigure{semi_anal}"
    },
    "0002/astro-ph0002460_arXiv.txt": {
        "abstract": "We present a study of the hard X-ray properties of a sample that includes all the Luminous Infrared Galaxies (LIGs, $\\rm L_{IR} > 10^{11}$ L$_\\odot$) observed in the 2-10 keV energy band (new and archival data). We find that a significant fraction of the sources optically classified as AGNs do not show any indication of nuclear activity in the X rays, thus suggesting heavy absorption along our line of sight. The absence of strong emission in the 20-200 keV band in a subsample of LIGs observed with BeppoSAX suggests that in many cases these sources are completely Compton thick ($\\rm N_H > 10^{25}$ cm$^{-2}$). From a comparison between the infrared and the X-ray emission we deduce that the mid-IR emission is absorbed by a lower column density than the X-ray emission or, alternatively, that the dust-to-gas ratio is lower than Galactic. We describe a simple model that reproduces the IR--X correlation by means of mixed AGN and starburst contributions and we compare the predictions of this model with the observational data at X-ray and optical wavelengths. Finally, we discuss the biases that affect the currently available samples of LIGs and briefly analyze a small unbiased sample, finding that at least 50\\% of the sources host a (weak) AGN. ",
        "introduction": "The nature and the energy source of the Luminous Infrared Galaxies (LIGs) has been a matter of debate since their discovery by IRAS, more than 10 years ago. This class of sources is characterized by a very high luminosity emitted mostly in the far--infrared ($\\rm L_{IR} > 10^{11} L_{\\odot}$). Such a luminosity is higher than in normal galaxies ($\\rm L_{IR}\\sim 10^{10}$ erg s$^{-1}$ for a large spiral like the Milky Way or M31) and therefore an additional source of energy is required. The infrared emission of LIGs is due to the presence of a large amount of dust that reprocesses optical and UV radiation to IR wavelengths. Since the primary emission is often not directly observable, its origin remains controversial.  Two main mechanisms have been invoked in this context: an intense starburst activity and/or an Active Galactic Nucleus (AGN). It is now widely accepted that in general both sources can be present in the LIGs, but the relative contribution of each is still unclear.  Recent studies with ISO of a sample of ULIGs (Ultra Luminous Infrared Galaxies, $\\rm L_{IR} > 10^{12} L_{\\odot}$) indicate that the bulk of the IR emission is due to starburst activity (Genzel et al. 1998, Lutz et al. 1998). The diagnostic used in these works is mainly based on mid-infrared spectroscopy: in particular, the policyclic aromatic hydrocarbon features at 7.7 $\\mu$m are claimed to be better than optical lines at discriminating between starburst and AGN, because of the lower extinction.  On the other side, numerous X-ray observations performed in the last few years, which will be in part reviewed in this paper, suggest that an AGN contribution to the observed IR luminosity is common.  Resolving this issue would be important not only for the understanding of the LIG phenomenon, but also to assess the contribution of the LIGs to the diffuse X-ray background (XRB). The knowledge of the fraction of the bolometric luminosity which emerges in the hard X-ray domain would allow the inclusion of the LIGs in the models of the XRB. These models are by and large successful in synthesizing the XRB from the contributions of individual AGNs; they must however include a dominant contribution from absorbed, type 2 AGNs, which up to now have been observed directly only at low luminosities and low redshifts. The zeroth-order extrapolation, which assumes that type 2s and type 1s evolve in the same way (``unified model''), has to face several discrepancies, and could be cured only by adding extra type 2 sources at intermediate or high redshifts (Gilli et al. 1999). Some of the LIGs could very well be the required additional sources, but their mean X-ray properties are still poorly known.  The expected X-ray spectrum, and the relative contribution of the X-ray emission to the bolometric luminosity are very different in the cases of AGN and starburst dominance, therefore hard (2-10 keV) X-ray data can be a powerful diagnostic.  The typical X-ray starburst spectrum is a thermal continuum with temperatures ranging from a fraction of a keV to several keV (in units of kT), plus several emission lines at low energies, and a K$_{\\alpha}$ feature of highly ionized iron at the energy of E$\\simeq 6.7$ keV (Cappi et al. 1999, Persic et al. 1998). The 2-10 keV luminosity is less than one part in 10$^3$ of the total emission.  The AGN spectrum between 2 and 10 keV is strongly dependent on the amount of obscuration suffered by the primary continuum component: if the absorbing column density $\\rm N_H$ is lower than $10^{24}$ cm$^{-2}$, the direct continuum spectrum is visible at energies higher than the photoelectric cutoff E$_C$, whose value depends on N$_H$. The mean direct continuum, as deduced from the spectra of unabsorbed AGNs, is well represented by a powerlaw with photon index $\\Gamma\\sim 1.7$. In addition, a K$_{\\alpha}$ iron line is observed at energies ranging from 6.4 keV (neutral iron) to 6.95 keV (hydrogen-like iron). The current view is that the iron line originates both from the accretion disc and from a reprocessing medium located farther away (1--10 pc): the first component can be broadened by relativistic effects and is suppressed if the absorbing column density is higher than several times 10$^{23}$ cm$^{-2}$, the second component is narrow, and its energy depends on the ionization degree of the reprocessing medium. The equivalent width (EW) of the line is EW $\\sim$ 100-200 eV with respect to the intrinsic continuum, and can be much higher (up to some keV) if the continuum is heavily absorbed at the line energy (while the line, produced in a different region, is still visible).  If the column density is higher than $10^{24}$ cm$^{-2} $ (i.e. the medium is Compton thick), the intrinsic emission is completely absorbed and only a reflected component survives. The X-ray spectrum of Compton thick sources is generally characterized by an Fe line with large equivalent width ($\\sim$ 1 keV) and a reflection dominated continuum flatter than the intrinsic spectrum (see Maiolino et al. 1998 for further details). While the {\\it observed}~2--10 keV contribution to the bolometric luminosity can be higher than 10\\% in unabsorbed AGNs, in Compton thick Seyfert 2s it can be lower than 0.1\\%.  Summarizing, the X-ray continuum of a starburst is expected to have a thermal high energy cutoff  and a soft spectrum with emission lines, while an AGN is characterized by a hard power law with a possible low energy cutoff due to absorption or, in the extremely absorbed cases (Compton thick), a flat reflection spectrum. Cold (E$\\sim$ 6.4 keV) and broad iron lines, or a large relative X-ray luminosity are clear signatures of an AGN, whereas narrow lines of ionized iron or a relatively weak X-ray emission are possible in both scenarios.  In this paper we study the sample of all the luminous infrared galaxies observed so far in hard X rays (2-10 keV). We combine the X-ray information with the IRAS photometry, and find interesting correlations between X-ray properties and IR colours. We propose a simple model that matches the observed correlations and is compatible with the available spectroscopic and spectropolarimetric data at near-IR and optical wavelengths.  In Sect. 2 we describe in detail our sample, in Sect. 3 we analyze the X-ray properties of the sources and the correlation between the X and IR emission. Finally, in Sect. 4 we discuss our results in the framework of a geometrical two-parameter model. Conclusions and future work are summarized in Sect. 5. In the Appendix we briefly report the results of our own spectral analysis of those sources which have not been published elsewhere. Throughout this paper we use the cosmological parameters H$_0=75$ km s$^{-1}$ Mpc$^{-1}$ and q$_0=0.5$. ",
        "conclusions": "This paper deals with the sample of luminous infrared galaxies (LIGs) observed to date in hard X~rays (2-10 keV). The sample is affected by a selection bias in favor of AGN-dominated sources, nevertheless it covers a region of the parameter space large enough to allow some general conclusions. For a significant fraction of the sample (15/78) hard X-ray data are analyzed and published for the first time.  The main results of our work can be summarized as follows: \\medskip  \\noindent 1) The X-ray properties of the LIGs are very different from source to source: the X-ray brightest objects are Seyfert~1 and low--absorption Seyfert~2 galaxies, with a X/IR flux ratio as high as 0.1-0.2. At the opposite extreme are the starburst dominated objects, for which the X/IR flux ratio is $\\sim 10^{-4}$ or lower. Interestingly, we note that a significant fraction of the sources optically classified as AGNs do not show any indication of nuclear activity in X~rays. Thus, either the AGN contribution is negligible, or the direct emission must be absorbed by a column density $\\rm N_H > 10^{24}$ cm$^{-2}$. \\medskip  \\noindent 2) A subsample of 13 LIGs was observed up to 200 keV. Six out of 11 sources that are Compton--thick in the 2-10 keV range are reflection dominated also from 10 to 200 keV, thus implying a column density $\\rm N_H>10^{25}$ cm$^{-2}$ (or, alternatively, a fading nucleus). Only two have an excess in the 15-100 keV range, while for the last three the hard 15-100 keV X-ray emission remains unconstrained. This result is somewhat puzzling, for the 10$^{24} - 10^{25}$ cm$^{-2}$ interval of column density appears to be underpopulated both for the LIGs and the non-LIG local Seyfert~2 galaxies. \\medskip  \\noindent 3) The shape of the correlation between the X/IR flux ratio and the 25/60 $\\mu$m IR colour suggests that the 25$\\mu$m emission is absorbed by a lower column density than the X-ray emission or, alternatively, that the absorbing material has a dust--to--gas ratio lower than Galactic.  The X-IR correlation is well reproduced by a model in which both AGNs and starbursts contribute to the total emission. The differences in the IR colours are mainly due to the different contribution of the starburst component, while the X/IR flux ratio is mainly determined by the amount of absorption affecting the AGN.  Additional evidence in support of the model is the detection of broad optical polarized or infrared lines only in the IR-warmest sources (where we see directly the inner region of the torus), and the steepness of the X-ray spectra of the IR-colder objects (where the starburst contributes strongly). \\medskip  \\noindent 4) Finally, we assembled an unbiased sample of LIGs from the BGS catalog, in a way which maximizes the fraction of sources with hard X-ray data available (15 out of 23, i.e. 65\\%). The sample is limited in {\\it total} IRAS flux (i.e. from 12$\\mu$m to 100$\\mu$m). This selection criterion removes a bias against AGNs present in most other LIG samples that, instead, are limited in the 60$\\mu$m flux. From the optical classification and from the IR and X-ray data we find that AGNs are common among the LIGs (13 out of 23 host an active nucleus) but they are either weak or heavily obscured."
    },
    "0002/astro-ph0002474_arXiv.txt": {
        "abstract": "We present Keck R and I band images of the field of the anomalous \\Xray\\ pulsar \\bron.  We derive an improved X-ray position from archival ROSAT HRI observations by correcting for systematic (boresight) errors.  Within the corresponding error circle, no object is found on the Keck images, down to limiting magnitudes $R=25.7$ and $I=24.3$.  We discuss the constraints imposed by these limits, and conclude that it is unlikely that \\bron\\ is powered by accretion from a disk, irrespective of whether it is in a binary or not, unless the binary is extremely compact. ",
        "introduction": "\\label{intro}  There is a growing number of \\Xray\\ pulsars for which it is not clear why they emit \\Xray{}s.  Their rotational energy loss appears insufficient to power the \\Xray\\ luminosity and there is no sign of a companion from which they could accrete matter.  These pulsars are called the ``anomalous X-ray pulsars'' (AXP).  The AXP differ from other X-ray pulsars in the following properties (\\cite{Hellier}, \\cite{MS95}, \\cite{Paradijs}): (i) their spin periods are within a small range, between 5 and 9 seconds, and their spin-down rates are nearly constant; (ii) timing analysis gives no evidence of orbital motion; (iii) their \\Xray\\ spectra are rather soft; (iv) their \\Xray\\ light curves show less variability; (v) their inferred \\Xray\\ luminosities are rather low ($\\sim\\!10^{35}\\,\\mathrm{erg\\,s^{-1}}$); (vi) they are located close to the Galactic plane; (vii) some appear associated with supernova remnants.  Members of the group of AXP include \\object{4U 0142+61}, \\bron, \\object{1E 1048.1$-$5937} and \\object{RX J1838.4$-$0301} (\\cite{MS95} and references therein).  Recently, three \\Xray\\ pulsars with similar properties were discovered: \\object{1RXS J170849.0$-$400910} (\\cite{Sugizaki}); \\object{1E 1841$-$045} in the supernova remnant \\object{Kes 73} (\\cite{Kes73}); and \\object{AX J1845.0$-$0300} (\\cite{Torii}; \\cite{GV}) in the supernova remnant G29.6+0.1 (\\cite{GGV99}). Indeed, perhaps even the progeny of these systems has been identified, in \\object{RX J0720.4$-$3125} (\\cite{Haberl}).  The latter object is the only one for which the optical counterpart has been identified (\\cite{MH98}; \\cite{KvK98}).  A variety of models for AXPs has been put forward, using both accretion and intrinsic mechanisms to explain the \\Xray\\ luminosity. For the accretion scenarios, the source of matter has been suggested to be a very low mass companion (\\cite{MS95}) or a circumstellar disk composed of debris, either from a common envelope phase of a high mass \\Xray\\ binary (\\cite{Paradijs}; \\cite{Ghosh}) or from a supernova explosion (\\cite{Corbet}; \\cite{Chatterjee}).  In scenarios in which the \\Xray\\ emission is intrinsic, the sources are either hot, massive, rapidly spinning white dwarfs, formed in the merger of a double white-dwarf binary (\\cite{Paczynski}), or so-called ``magnetars,'' isolated neutron stars with very high magnetic field strengths ($\\ga\\!10^{14}\\,\\mathrm{G}$; \\cite{TD96}).  In the latter case, AXP might be related to soft gamma-ray repeaters\\footnote{The proposed \\Xray\\ counterparts to SGRs also have spin periods in the 5--9 s range (\\cite{SGR1806}, \\ctyr{SGR1900}).}.  In this paper, we present a deep search for the optical counterpart to the AXP \\bron.  In Sect.~\\ref{sec:bron}, we briefly review what is known about this source.  We derive a refined position for \\bron\\ in Sect.~\\ref{sec:position}, and in Sect.~\\ref{sec:keck} we present our optical observations.  We discuss the results in Sect.~\\ref{sec:discuss}. ",
        "conclusions": "We have shown that the tight limits we derive on the optical emission of \\bron\\ are a serious problem for scenarios involving accretion by way of an accretion disk.  This is irrespective of whether the source is in a binary, except if the binary is extremely compact. Therefore, we conclude that, most likely, \\bron\\ is an isolated object, and that its emission is not powered by accretion. We note that the same conclusion was indepently drawn by Kaspi et~al. (\\ctyr{Kaspi}) on the basis of the extreme rotational stability of \\bron."
    },
    "0002/astro-ph0002197_arXiv.txt": {
        "abstract": "We present the first spectra obtained with the {\\em Very Large Telescope} for white dwarfs in a globular cluster. Estimates of atmospheric parameters are obtained and compared to evolutionary tracks. We discuss possible implications for the distance scale of globular clusters and white dwarf evolution and demonstrate how white dwarfs might be used to establish an independent distance scale to globular clusters. ",
        "introduction": "White dwarfs are the final stage of all low-mass stars and therefore all single stars in a globular cluster that currently finish their nuclear-burning lifetimes are expected to evolve into white dwarfs. As this has been the situation for many billions of years globular clusters should contain many white dwarfs. However, these stars managed to evade detection until photometric white dwarf sequences in globular clusters were discovered recently by observations with the {\\em Hubble Space Telescope} (HST) (Paresce et al.\\ \\cite{pade95}, Richer et al.\\ \\cite{rifa95}, \\cite{rifa97}, Cool et al.\\ \\cite{copi96}, Renzini et al.\\ \\cite{rebr96}). Photometric observations contain only a limited amount of information: The two chemically distinct white dwarfs sequences (hydrogen-rich DA's and helium-rich DB's) in principle can be distinguished by their photometric properties alone in the temperature range $10,000\\,K \\leq T_{\\mathrm{eff}} \\leq 15,000$\\,K (see Bergeron et al.\\ \\cite{bewe95}). Renzini et al.\\ (\\cite{rebr96}) classified two white dwarfs in NGC\\,6752 as DB's by this method. However, without a spectral classification, both stars can also be explained as high mass DA white dwarfs, possibly a product of merging. Richer et al.\\ (\\cite{rifa97}) speculate that the brightest white dwarf in M\\,4 (V=22.08) might be a hot (27,000K) DB star. \\newline\\indent The location of the white dwarf cooling sequence (and thus the brightness of the white dwarfs) is also sensitive to the white dwarf mass. Renzini et al.\\ (\\cite{rebr96}) argued that the white dwarf masses in globular clusters are constrained to the narrow range 0.51\\Msolar $\\leq {\\rm M_{WD}} \\leq$ 0.55\\Msolar, but some systematic differences between clusters are obvious: At a given metallicity some globular clusters (e.g.\\ NGC\\,6752) possess very blue horizontal branches (HB's) with HB star masses as low as 0.50\\Msolar. Such extreme HB stars evolve directly to low-mass C/O white dwarfs (bypassing the AGB), shifting the mean white dwarf mass closer to 0.51\\Msolar. Other clusters show only red HB stars, which will evolve to the AGB and form preferably white dwarfs with masses of $\\approx$0.55\\Msolar. \\begin{figure} \\vspace{8cm} \\special{psfile=b1293.f1 hscale=100 vscale=100 hoffset=-130 voffset=-330 angle=0} \\caption{Colour-magnitude diagram of NGC~6397 (King et al.\\ \\cite{kian98}, their Fig.~2). Open circles mark the four white dwarfs for which spectra could be obtained, the open square marks WF4-205 (see text).\\label{cmd}} \\end{figure} Low-mass white dwarfs (M$<$0.45\\Msolar) with a degenerate He core are produced if the red giant branch evolution is terminated by binary interaction before the helium core exceeds the minimum mass for helium burning. Recently, Cool et al.\\ (\\cite{cogr98}) found 3 faint UV-bright stars in NGC\\,6397 which they suggest could be He white dwarfs (supported by Edmonds et al.\\ \\cite{edgr99}). Massive white dwarfs may be produced from blue stragglers or by collisions of white dwarf-binaries with subsequent merging (e.g. Marsh et al.\\ \\cite{madh95}). \\newline\\indent Only a detailed spectroscopic investigation can provide masses and absolute luminosities of the individual globular cluster white dwarfs. This is also very important for the use of white dwarfs as standard candles to derive distances to globular clusters (Renzini et al.\\ \\cite{rebr96}): The basic idea is to fit the white dwarf cooling sequence of a globular cluster to an appropriate empirical cooling sequence of local white dwarfs with well determined trigonometric parallaxes. The procedure is analogous to the classical main sequence fitting but avoids the complications with metallicity -- white dwarfs have virtually metal free atmospheres. In addition they are locally much more abundant than metal-poor subdwarfs. The arrival of the {\\sc Hipparcos} results as well as new metallicity determinations have rekindled the debate on globular cluster distances (see the review by Reid \\cite{reid99} and references therein). A further check on the distance is therefore urgently needed. \\newline\\indent We started an observing programme at the ESO {\\em Very Large Telescope (VLT)} to obtain spectra of white dwarfs in globular clusters. The programme consists of two parts: First, low S/N ($\\approx 10$) spectra of the white dwarf candidates are obtained to verify their spectral type and estimate their effective temperatures. In a second run we plan to observe higher S/N ($\\approx 30$) spectra that will allow to derive \\logg\\ with an internal error of $\\le 0.1$dex. Here we report on the very first results for NGC~6397.  \\begin{table} \\caption{Target coordinates and photometric data (Cool priv. comm.). \\label{targ}} \\begin{tabular}{lrrrr} \\hline Star & $\\alpha_{2000}$ & $\\delta_{2000}$ & $V$ & $V-I$\\\\ \\hline WF4-358 & \\RA{17}{40}{58}{52} & \\DEC{$-$53}{42}{25.3} & \\magpt{22}{73} & \\magpt{+0}{02}\\\\ WF2-479 & \\RA{17}{41}{07}{06} & \\DEC{$-$53}{44}{45.1} & \\magpt{23}{87} & \\magpt{+0}{22}\\\\ WF4-205 & \\RA{17}{41}{01}{58} & \\DEC{$-$53}{43}{15.3} & \\magpt{23}{99} & \\magpt{+0}{30}\\\\ WF2-51 & \\RA{17}{41}{04}{52} & \\DEC{$-$53}{43}{55.8} & \\magpt{24}{00} & \\magpt{+0}{28}\\\\ WF2-846 & \\RA{17}{41}{01}{78} & \\DEC{$-$53}{44}{47.2} & \\magpt{24}{30} & \\magpt{+0}{33}\\\\ \\hline \\end{tabular} \\end{table} \\begin{figure} \\vspace{8.5cm} \\special{psfile=b1293.f2 hscale=43 vscale=43 hoffset=-10 voffset=-70 angle=0} \\caption{Traces through the slits along the spatial axis. The lines mark the white dwarf spectra (WF4-205 unfortunately lies on the wing of a much brighter star).\\label{slits}} \\end{figure}  \\begin{figure} \\vspace{8.5cm} \\special{psfile=b1293.f3 hscale=43 vscale=43 hoffset=-10 voffset=-70 angle=0} \\caption{The (relatively flux-calibrated) spectra of the white dwarfs in NGC~6397. The spectra of the three faintest stars are offset from each other as they would otherwise overlap. WF4-358 has no additional offset relative to WF2-479.\\label{spectra}} \\end{figure} ",
        "conclusions": "Using VLT-FORS1 multi object spectroscopy we have confirmed four white dwarf candidates to be hydrogen-rich DA white dwarfs. The gravity determined for the brightest star, WF4-358, suggests that it could be a He white dwarf with a mass of (0.36$\\pm$0.12)\\Msolar, although the error bars are large enough to also accommodate a C/O white dwarf. Temperatures derived for the three cooler and fainter stars for fixed \\logg\\ would put them near the red edge of the ZZ Ceti instability strip, for which Bergeron et al.\\ (\\cite{bewe95b}) determine a temperature range of 11,160~K to 12,460~K using their preferred ML2/$\\alpha$=0.6 prescription for the treatment of convection. Therefore, a search for photometric variability of WF2-479 and WF2-51 -- if successful -- could place important additional constraints on these stars. The systematic variation of the distance moduli derived for a given \\logg\\ shows that the assumption of a constant mass for all white dwarfs in a globular cluster may bias a distance determination. However, due to the low S/N of the current data, these results are preliminary. Once higher quality spectra are available, which will allow more accurate parameter (and thus mass) determinations, analyses of white dwarfs in globular clusters will become a powerful tool for independent distance estimates."
    },
    "0002/astro-ph0002368_arXiv.txt": {
        "abstract": "The relation between quasar variability and parameters such as luminosity and redshift has been a matter of hot debate over the last few years with many papers on the subject.  Any correlations which can be established will have a profound effect on models of quasar structure and evolution.  The sample of quasars with redshifts in ESO/SERC field 287 contains over 600 quasars in the range $0 < z < 3.5$ and is now large enough to bin in luminosity and redshift, and give definitive measures of the correlations.  We find no significant correlation between amplitude and redshift, except perhaps at very low redshift, but an inverse correlation between amplitude and luminosity.  This is examined in the context of various models for quasar variability. ",
        "introduction": "One of the most important constraints on the structure of quasars is variability.  Short term variations set limits on the size of the emitting region, and differences in the nature of the variation in the Xray, optical and radio domains give clues to the underlying structure. Variations on longer timescales of a few years are not so easily accounted for by the current black hole paradigm, and microlensing has been put forward as an alternative explanation for the observations. Although the most extensive monitoring of quasar variability has been done in optical wavebands, there is still little consensus even over the broad picture.  One of the first problems is to parametrise the variation in such a way that it can be compared with models, or even more general expectations.  The timescale of variation and the amplitude are the two parameters which have been mostly studied, although some authors have succeeded in confusing the two.  This typically involves claiming that in a short run of data, objects varying on a short timescale will achieve a larger amplitude than those varying on a longer timescale, and so amplitude can be taken as an (inverse) measure of timescale (Hook et al. \\cite{hoo94}).\\\\  In this paper we shall confine our attention to the correlation of amplitude with other parameters.  Amplitude is an easy parameter to measure, and involves none of the difficulties inherent in estimating timescale of variation.  There are modes of variability in which some uncertainty will be introduced by the length of the run of observations, but this is a problem which can arise in any time series analysis. Most attention has so far been given to possible correlations of amplitude with redshift or luminosity (eg Trevese et al. \\cite{tre89}, Giallongo et al. \\cite{gia91}, Hook et al. \\cite{hoo94}, Cristiani et al. \\cite{cri96}), but even in such straightforward situations there has been little agreement.\\\\  The reasons for the lack of consensus are not easy to understand, but it is clear that any effect which is present is not large compared with the cosmic scatter in the data, which in the case of amplitude appears to be about one magnitude, much larger than any photometric errors.  Perhaps the most likely cause of disagreement is in the selection of samples for analysis.  In order to obtain a meaningful measure of correlation it is essential that unbiassed samples are used, and that they cover a large range of redshift and luminosity.  In this paper particular attention is paid to the sample selection process, and by extending the analysis to high redshifts the baseline for the measurement of correlations is greatly extended over earlier samples.\\\\ ",
        "conclusions": "In this paper we have examined the dependance of amplitude on redshift and luminosity for large samples of quasars selected on the basis of ultra-violet excess and variability.  The quasars span a redshift range $0 < z < 3.5$ and a luminosity range $-20 > M_{B} > -28$. There is evidence for a correlation of amplitude with luminosity and/or redshift for the sample as a whole, but when it is binned in redshift the correlation with luminosity becomes significant.  This result could be strengthened by the possible existence of a population of large amplitude low luminosity objects in the sample.  No convincing evidence is found for a correlation between amplitude and redshift, either for the sample as a whole or when it is binned in luminosity.\\\\  Various models of quasar variability are examined with respect to the observed correlations.  It is concluded that any straightforward interpretation of an accretion disk model is incompatible with the data.  The Christmas Tree model has some merits, especially in the regime of low luminosity quasars and Seyfert galaxies, but for luminous quasars the rate of supernovae required is too large to be compatible with the observed variability.  The microlensing model can be used to explain all the data, although it does require that the Einstein radius of the microlensing bodies is comparable in size to the quasar nucleus.\\\\"
    },
    "0002/astro-ph0002018_arXiv.txt": {
        "abstract": "In order to explore how the FR~I / FR~II dichotomy is related to the nuclear properties of radio galaxies, we studied a complete sample of 26 nearby FR~II radio galaxies using Hubble Space Telescope (HST) images and compared them with a sample of FR~I previously analyzed.  FR~I nuclei lie in the radio-optical luminosity plane along a tight linear correlation, which argues for a common synchrotron origin.  FR~II show a more complex behavior, which is however clearly related to their optical spectral classification.  Broad line FR~II radio galaxies (BLRG) are located overall well above the FR~I correlation, suggesting that a contribution from thermal (disc) emission is present. Three narrow line (NLRG) and one weak line radio galaxy (WLRG), in which no nuclear source is seen, can be interpreted as the obscured counter-parts of BLRG, in agreement with the current unification schemes.  Conversely, in 5 sources of the sample, all of them NLRG or WLRG, optical cores are located on the same correlation defined by FR~I and with similar radio and optical luminosities.  This suggests that, in analogy to FR~I, the emission is dominated by synchrotron radiation and represents the optical counter-part of the non--thermal radio cores. Interestingly, all these galaxies are located in clusters, an environment typical of FR~I.  These results imply that, at least at low redshifts, the FR~II population is not homogeneous. Furthermore, the traditional dichotomy between edge darkened and brightened radio morphology is not univocally connected with the innermost nuclear structure, as we find FR~II with FR~I--like nuclei and this has interesting bearings from the point of view of the AGN unified models. ",
        "introduction": "\\label{intro} The original classification of extended radio galaxies by Fanaroff \\& Riley (\\cite{fr}) is based on a morphological criterion, i.e. edge darkened (FR~I) vs edge brightened (FR~II) radio structure.  It was later discovered that this dichotomy corresponds to a (continuous) transition in total radio luminosity (at 178 MHz) which formally occurs at $L_{178}= 2\\times 10^{33}$ erg s$^{-1}$ Hz$^{-1}$.  The presence of radio sources with intermediate morphology, in which typical FR~I structures (such as extended plumes and tails) are seen together with features characteristics of FR~II sources (narrow jets and hot spots) (see e.g. Parma et al.  \\cite{parm87}, Capetti et al. \\cite{capetti95}) argues in favour of a continuity between the two classes.  From the optical point of view both FR~I and FR~II are associated with various sub-classes of elliptical-like galaxies, but statistically their populations are different (Zirbel \\cite{zirb96}).  Owen (\\cite{owen93}) found that the FR~I/FR~II division is also linked to the optical magnitude of the host galaxy, possibly suggesting that the environment plays an important role in producing different extended radio morphologies. Moreover, FR~II are generally found in regions of lower galaxy density and are more often associated with galaxy interactions with respect to FR~I (Prestage \\& Peacock \\cite{pp}, Zirbel \\cite{zirb97}). Differences are also observed in the optical spectra: while FR~I are generally classified as weak-lined radio galaxies, strong (narrow and broad) emission lines are often found in FR~II (Morganti et al. \\cite{morg92}, Zirbel \\& Baum \\cite{zirb95}), although a sub-class of weak-lined FR~II is also present (Hine \\& Longair \\cite{hine}).  Within the unification scheme for radio-loud AGN (for a review, see Urry \\& Padovani \\cite{urry95}), FR~I and FR~II radio galaxies are thought to represent the parent population of BL Lac objects and radio-loud quasars, respectively (Antonucci \\& Ulvestad \\cite{anto85}, Barthel \\cite{bart89}).  In order to explain the lack of broad lines in the ``mis-oriented'' (narrow-lined) FR~II--type objects, obscuration by a thick torus is invoked. A combination of obscuration and beaming is therefore necessary at least for the FR~II-quasars unification (e.g. Antonucci \\& Barvainis \\cite{anto90}). However, there is evidence that this simple picture is probably inadequate: some radio--selected BL Lacs - among the most powerful sources in the class - display an extended radio structure and luminosity typical of FR~II (Kollgaard et al. \\cite{koll92}, Murphy et al. \\cite{murphy93}) and broad - although weak - lines have been observed in some BL Lacs. Moreover, Owen et al. \\cite{owen96} noted that the lack of BL Lacs in a sample of radio galaxies located in Abell clusters can be an effect of their selection criteria if the parent population of BL Lacs includes both FR~I and FR~II. This idea is also consistent with a recently proposed modification of the unification scheme, which claim that the weak-lined FR~II are indeed associated with BL Lac objects (Jackson \\& Wall \\cite{jackson}). These observations can be however reconciled with the unification scenario once continuity between the weak and powerful radio--loud sources is allowed and thus transition objects are expected.  In Chiaberge et al. (\\cite{pap1}, hereafter Paper~I) we studied HST images of all FR~I radio galaxies belonging to the 3CR catalogue, finding that unresolved nuclear sources are commonly present in these objects. A strong linear correlation is found between this optical and the radio core emission, extending over four orders of magnitude in luminosity. This, together with spectral information, strongly argues for a common non-thermal origin, and suggests that the optical cores can be identified with synchrotron radiation produced in a relativistic jet, qualitatively supporting the unifying model for FR~I and BL Lacs. Furthermore, the high rate detection (\\gta 85 \\%) of optical cores in the complete sample indicates that a standard pc--scale geometrically thick torus is not present in these low-luminosity radio galaxies. Any absorption structure, if present, must be geometrically thin, and thus the lack of broad lines in FR~I cannot be attributed to obscuration. Alternatively, thick tori are present only in a minority of FR~I. Given the dominance of non-thermal emission, the optical core luminosity also represents a firm upper limit to any thermal component, suggesting that accretion might take place in a low efficiency radiative regime.  The picture which emerges from this analysis is that FR~Is lack substantial thermal (disc) emission, Broad Line Regions and obscuring tori, which are usually associated with radio-quiet and powerful radio--loud AGN.  As a natural extension of Paper~I, here we study the HST images of a sample of low redshift FR~II radio galaxies, in order to explore how the differences in radio morphology are related to the optical nuclear properties. In particular, one of the most important questions is whether the FR~I/FR~II dichotomy is generated by two different manifestations of the same astrophysical phenomenon, and the transition between the two classes is indeed continuous, or instead it reflects fundamental differences in the innermost structure of the central engine.  The selection of the sample is presented and discussed in Sect. \\ref{thesample}, while in Sect. \\ref{hstobs} we describe the HST observations.  In Sect. \\ref{CCCII} we focus on the detection and photometry of the optical cores. Finally, in Sect. \\ref{discussion} we discuss our findings. ",
        "conclusions": "In Chiaberge et al. (\\cite{pap1}) we discovered that FR~I nuclei lie in the radio-optical luminosity plane along a tight linear correlation.  We argued that this is due to a common synchrotron origin for both the radio and optical emission. FR~I nuclei must also be unobscured and intrinsically lacking of BLR and of significant thermal emission from any powerful accretion disc.  In order to explore how the differences in radio morphology are related to the optical nuclear properties, we analyzed HST images of 24 extended radio-galaxies morphologically classified as FR~II, belonging to the 3C catalog and with $z<0.1$. We detected optical cores in 13 sources, which implies that the covering fraction of any obscuring material is less than $\\sim 0.54$, or equivalently, the torus has an opening angle of $\\sim 63^{\\circ}$. This can be even larger if at least some of the upper limits are actually just below the detection threshold. Notice that our determination of this critical angle is inconsistent with the division between higher redshift ($0.5 < z < 1$) 3CR quasars and radio galaxies, which has been found to be $\\theta \\sim 45^{\\circ}$ (Barthel \\cite{bart89}). This might be a problem, however the low redshift selection of our sources does not allow to derive any firm conclusion. We are currently studying a larger and higher redshift sample in order to further investigate this issue (Chiaberge et al. in preparation).  Our results suggest that the radio morphology is not univocally connected with the optical properties of the innermost structure of radio galaxies.  In fact, at least at low redshifts, there is not a single homogeneous population of FR~II: unlike FR~I, they show a complex behavior, which is however clearly related to their optical spectral classification.  In BLRG optical nuclei are likely to be dominated by thermal (disc) emission. As discussed above, line emission contamination cannot account for this excess. In agreement with the current unification scheme of radio loud AGNs, we also identify their possible obscured counter--parts. It seems that broad lines and obscuring tori are closely linked and both are present only associated to radiatively efficient accretion.  We also find five FR~II sources, spectrally identified as narrow lined objects, which harbor nuclei essentially indistinguishable from those seen in FR~I. By analogy with FR~I, we argue that their optical nuclear emission is produced primarily by synchrotron radiation, they are not obscured to our line of sight and therefore intrinsically lack a BLR.  Clearly, a classification based on the optical nuclear properties, as seen in these HST images, is more likely to reflect true similarities (or differences) on the nature of the central engine (such as, e.g., the rate of radiative dissipation in the accretion disc) than the traditional dichotomy of radio morphology.  From our data and within the limits of the available statistics, we find no evidence of a continuous transition between the two classes (FR~I and FR~II), as they are well separated in the $L_r$ vs $L_o$ plane. At this stage we only point out that sources with cores below $L_o < 10^{27.5}$ erg s$^{-1}$ Hz$^{-1}$ (or equivalently $L_r < 10^{31}$ erg s$^{-1}$ Hz$^{-1}$) have FR~I like nuclei, while FR~II start above this threshold.  It is of particular interest that a significant fraction of FR~II (at least 30 \\%, but can be as large as 50 \\% depending on the nature of the sources without detected optical nuclei) have FR~I--like nuclei.  The fact that all of these are located in clusters, an environment typical of FR~I, might represent an important hint on the origin of the different flavours of radio galaxies, worth exploring through the study of a larger sample of objects.  These results have also interesting bearings from the point of view of the unified models. In fact, this picture argues against the idea that all FR~II radio galaxies constitute the parent population of radio-loud quasars. We propose instead that galaxies with FR~II morphology and an FR~I-like core are possibly mis-aligned counter--parts of BL Lac objects. This can account for the observation that some radio-selected-type BL Lacs show radio morphologies more consistent with FR~II than with FR~I (e.g. Kollgaard et al. \\cite{koll92}).  To conclude, we note that all of the galaxies included in our sample are low redshift objects with total radio powers not exceeding $L_{178} \\sim 10^{27}$ erg s$^{-1}$ Hz$^{-1}$: thus a crucial observational issue is to understand whether these results hold to higher power/redshift samples or they are limited to low luminosities FR~II. This will be explored in a forthcoming paper."
    },
    "0002/astro-ph0002062_arXiv.txt": {
        "abstract": "We present measurements of [Fe/H] for six M supergiant stars and three giant stars within 0.5 pc of the Galactic Center (GC) and one M supergiant star within 30 pc of the GC. The results are based on high-resolution ($\\lambda / \\Delta \\lambda =$ 40,000) $K-$band spectra, taken with CSHELL at the NASA Infrared Telescope Facility. We determine the iron abundance by detailed abundance analysis, performed with the spectral synthesis program MOOG. The mean [Fe/H] of the GC stars is determined to be near solar, [Fe/H] = +0.12 $\\pm$ 0.22. Our analysis is a {\\it differential} analysis, as we have observed and applied the same analysis technique to eleven cool, luminous stars in the solar neighborhood with similar temperatures and luminosities as the GC stars. The mean [Fe/H] of the solar neighborhood comparison stars, [Fe/H] = +0.03 $\\pm$ 0.16, is similar to that of the GC stars. The width of the GC [Fe/H] distribution is found to be narrower than the width of the [Fe/H] distribution of Baade's Window in the bulge but consistent with the width of the [Fe/H] distribution of giant and supergiant stars in the solar neighborhood. ",
        "introduction": "\\citet{mor96} and \\citet{ser96} have summarized the relative importance of star formation, gas inflow, and gas outflow in the Galactic Center (GC). Gas from the inner disk flows into the nucleus, perhaps driven by a bar potential \\citep{sta91,mor96}, and some of the stellar mass-loss from the bulge may fall into the GC \\citep{bli93,jen94}. \\citet{ser96} have proposed that this inflow of disk gas results in sustained star formation in a ``central molecular zone\" (CMZ) within a radius from the GC of $R \\leq$ 200 pc. The GC is observed to be currently forming stars in the CMZ at a rate of 0.3--0.6 M$_{\\odot}$ yr$^{-1}$ \\citep{gus89}. Some of the GC gas incorporated into stars, enriched by stellar nucleosynthesis, will be returned to the GC interstellar medium by stellar mass-loss and supernovae. X-ray observations of hot gas in the GC suggest that some gas could be driven temporarily or permanently from the GC by a galactic fountain or wind \\citep{bli93,mor96}, however, making it unclear how much enriched material is incorporated into following generations of star formation. The presence of strong magnetic fields in the GC also likely plays a role in whether enriched gas is driven from or is retained within the GC.  Continued star formation in the central few hundred parsecs of the Galaxy may lead to higher metallicities within the CMZ. Chemical abundances in the disks of spiral galaxies are observed to reach their highest values at the center \\citep[see ][]{pag81,shi90}. H II regions, planetary nebulae, and OB associations in the Milky Way also show a radial metallicity gradient \\citep{sha83,fic91,mac94, sim95,vil96,rud97,aff97,sma97,gum98}. Chemical evolution models strive to explain these gradients by considering the relative star formation and gas infall/outflow rates, and the metal abundance of the gas compared to the stars \\citep[see ][]{aud76,ran91}.  The metallicity of GC stars is needed in order to constrain models and understand the stellar processes in the central parsecs of our Galaxy. Measurements of stellar metallicities in the center of the Milky Way Galaxy, which is obscured by approximately 30 mag of extinction at $V$, are only now beginning to be feasible through infrared studies. \\citet{car00} found a solar Fe abundance in IRS 7, an M supergiant at a distance of $R = $ 0.2 pc from the GC, from a detailed abundance analysis of CSHELL ($\\lambda / \\Delta \\lambda $ = 40,000) $K-$band spectra. Our work analyzes similar data for several M supergiant and giant stars located in the central cluster, within 0.5 pc of the GC, and also one M supergiant star located in the Quintuplet cluster, 30 pc away from the GC. ",
        "conclusions": "We present the first measurement of stellar [Fe/H] for ten stars in the Galactic Center (GC), at distances from the GC of $R <$ 30 pc. Nine GC stars are located in the central cluster ($R <$ 0.5 pc) and one star is located in the Quintuplet cluster ($R$ = 30 pc). The abundance analysis is based on high-resolution ($\\lambda / \\Delta \\lambda =$ 40,000) $K-$band spectra. The mean [Fe/H] of the GC is determined to be near solar, [Fe/H] = +0.12 $\\pm$ 0.22, and also similar to the mean [Fe/H] for cool, luminous stars in the solar neighborhood, [Fe/H] = +0.03 $\\pm$ 0.16, observed and analyzed in the same way. The width of the GC [Fe/H] distribution, which ranges from [Fe/H] = --0.3 to +0.5, is found to be significantly narrower than the width of the [Fe/H] distribution of Baade's Window, which ranges from [Fe/H] = --1.0 to +0.8. The GC [Fe/H] distribution is consistent with the [Fe/H] distribution of supergiant stars in the solar neighborhood. This suggest that the most luminous stars in the GC are unlikely to be dominated by bulge-like stars, and that the evolutionary path of the GC, while unique, is closer to the disk's than to the bulge's. The Quintuplet star at $R$ = 30 pc has a similar [Fe/H] to stars located in the central cluster at $R <$ 0.5 pc. In the future, abundance measurements of CNO and $\\alpha-$elements are planned to provide a complete view of the abundance patterns of the stars in the central regions of the Milky Way."
    },
    "0002/astro-ph0002124_arXiv.txt": {
        "abstract": "s{Galaxy structures are certainly fractal up to a certain crossover scale $\\lambda_0$. A clear determination of such a scale is still missing. Usually the conceptual and practical implications of this property are neglected and the structures are only discussed in terms of their global amplitude. Here we present a compact summary of these implications. First, we discuss the problem of the identification of the crossover scale $\\lambda_0$ and the proper characterization of the scaling. We then consider the implications of these properties with respect to various physical phenomena and to the corresponding characteristic values, i.e.  $r_0$, $\\sigma_8$, $\\Omega$, etc.  These implications crucially depend on the value of $\\lambda_0$, but they are still important for a relatively small value, say $\\lambda_0 \\approx 50 \\hmp$. Finally we consider the main theoretical consequences of these results. } ",
        "introduction": "Nowadays there is a general agreement about the fact that galactic structures are fractal up to a distance scale of $\\sim 50 \\hmp$ \\cite{slmp98,jmsl99} and the increasing interest about the fractal versus homogeneous distribution of galaxy in the last year \\cite{coles98,scara98,rees99,cappi98,martinez99,hutton99,chown99,landy99,nl00} has focused, mainly on the determination of the homogeneity scale $\\lambda_0$.\\footnote{See the web page {\\it http://pil.phys.uniroma1.it/debate.html } where all these materials have been collected} The main point in this discussion is that galaxy structures are fractal no matter what is the crossover scale, and this fact has never been properly appreciated. Clearly, qualitatively different implications are related to different values of $\\lambda_0$.  \\begin{itemize}  \\item {\\it Characterization of scaling properties.}  Given a distribution of points, the first main question concerns the possibility of defining a physically meaningful average density. In fractal-like systems such a quantity depends on the size of the sample, and it does not represent  a reference value, as in the case of an homogeneous distribution. Basically a system cannot be homogeneous below the scale of the maximum void present in a given sample. However the complete statistical characterization of highly irregular structures is the objective of Fractal Geometry\\cite{man77}.   The major problem from the point of view of data analysis is to use statistical methods which are able to properly characterize scale invariant distributions, and hence which are also suitable to characterize an eventual crossover to homogeneity. Our main contribution \\cite{pie87,cp92,slmp98}, in this respect, has been to clarify that the usual statistical methods (correlation function, power spectrum, etc.) are based on the assumption of homogeneity and hence are not appropriate to test it. Instead, we have introduced and developed various statistical tools which are able to test whether a distribution is homogeneous or fractal, and to correctly characterize the scale-invariant properties. Such a discussion is clearly relevant also for the interpretation of the properties of artificial simulations. The agreement about the methods to be used for the analysis of future surveys such as the Sloan Digital Sky Survey (SDSS) and the two degrees Fields (2dF) is clearly a fundamental issue.  Then, if and only if the average density is found to be not sample-size dependent, one may study the statistical properties of the fluctuations with respect to the average density itself. In this second case one can study basically two different length scales. The first one is the homogeneity scale ($\\lambda_0$), which  defines the scale beyond which the density fluctuations become to have a small amplitude with respect to the average density ($\\delta \\rho < \\rho$). The second scale is related to the typical length scale of the structures of the density fluctuations, and, according to the terminology used in statistical mechanics \\cite{perezmercader}, it is called correlation length $r_c$. Such a scale has nothing to do with the so-called \"correlation length\" used in cosmology and corresponding to the scale $\\xi(r_0)=1$\\cite{pee80}, which is instead related to $\\lambda_0$ if such a scale exists.  \\item {\\it Implication of the fractal structure up to scale $\\lambda_0$}.  The fact that galactic structures are fractal, no matter what is the homogeneity scale $\\lambda_0$, has deep implication on the interpretation of several phenomena such as the luminosity bias, the mismatch galaxy-cluster, the determination of the average density, the separation of linear and non-linear scales, etc., and on the theoretical concepts used to study such properties. We discuss in  detail some of these points. We then review some of the main consequences of the power law behavior of the galaxy number density, by relating various observational quantities (e.g. $r_0$, $\\sigma_8$, $\\Omega$, etc.) to the length scale $\\lambda_0$.  We also note that the properties of dark matter are inferred from the ones of visible matter, and hence they are closely related. If now one observes different statistical properties for galaxies and clusters, this necessarily implies a change of perspective on the properties of dark matter.   \\item {\\it Determination of the homogeneity scale $\\lambda_0$. }  This is, clearly, a very important point which is at the basis of the understanding of galaxy structures and more generally of the cosmological problem. We distinguish here two different approaches: direct tests and indirect tests. By direct tests, we mean the determination of the conditional average density in three dimensional surveys, while with indirect tests we refer to other possible analyses, such as the interpretation of angular surveys, the number counts as a function of magnitude or of distance  or, in general, the study of non-average quantities, i.e. when the fractal dimension is estimated without making an average over different observes (or volumes). While in the first case one is able to have a clear and unambiguous answer from the data, in the second one is only able to make some weaker claims about the compatibility of the data with a fractal or a homogeneous distribution. For example   the papers of Wu et al\\cite{rees99} and  Nusser \\& Lahav\\cite{nl00} mainly concern with compatibility arguments, rather than with direct tests. However, also in this second case, it is possible to understand some important properties of the data, and to clarify the role and the limits of some underlying assumptions which are often used without a critical perspective. We do not enter here in the details of the discussion about real data (see e.g. Joyce et al. 1999, Wu et al. 1999), however in the last section we consider separately the case (i) $\\lambda_0 \\approx 50 \\hmp$, (ii) $\\lambda_0 \\approx 300 \\hmp$ and (iii) $\\lambda_0 \\approx 1000 \\hmp$, briefly discussing the main theoretical consequences. \\end{itemize} ",
        "conclusions": "We now present a short discussion about the perspective of our work: the cosmological implications of the fractal behavior of visible matter crucially depend on the crossover scale $\\lambda_0$, but, no matter what is the actual value of such a scale, we  have some important consequences from the theoretical point of view. We may identify three different scenarios.  \\begin{itemize}  \\item (1) The fractal extends only up to $\\sim 30\\div 100 \\hmp$. This is the minimal concept which begins to be absorbed in the literature\\cite{rees99} but, sometimes, without considering its real consequences. The standard approach to galaxy distribution has identified very small \"correlation lengths\", namely $5 \\hmp$ for galaxies and $25 \\hmp$ for clusters. These numbers (which were supposed to know with high precision\\cite{dav97}) are anyhow inconsistent with the fractal extending a factor $2\\div 4$ more. We have shown that this inconsistency is conceptual and not due to incomplete data or week statistics. Hence, in this hypothesis one has to abandon all the concepts related to these length scales. These are:  (i) The estimate of the matter density in clusterised objects (visible + dark), which has been claimed to be $\\Omega \\approx 0.2 \\div 0.3$, decreases by one order of magnitude or more.  (ii) The normalization of N-body simulations is usually performed to some length-scale or amplitude of fluctuations, which are related to $5 \\hmp$ and $25 \\hmp$.  (iii) Concepts like the galaxy-cluster mismatch and the related luminosity bias, as well as the understanding of the clustering via the bias parameter $b$ (i.e. linear or non-linear - \"stochastic bias\" - amplification of $\\xi(r)$) loose any physical meaning.  (iv) The interpretation of the velocity field is also based on the linear approximation which cannot certainly hold at scales smaller than $30\\div50 \\hmp$.  (v) The reconstruction of the three dimensional properties from the  angular data suffers of the same untested assumption of homogeneity.  In summary major modifications are necessary for the origin and dynamics of large scale structures and for the role  of dark matter. However the structures may still formed via gravitational instability, in the sense that they are not necessarily primeval\\cite{jampsl00}.  \\item (2) The fractal extends up to $300 \\div 500 \\hmp$. In this case the standard picture of gravitationally induced structures after the electromagnetic decoupling is untenable. There is no time to create such large scale correlated structures via gravitational instability, starting from Gaussian initial conditions. More string consequences are clearly important for what concerns the amount of matter in clusterised objects.  \\item (3) The fractal extends up to $\\gtapprox 1000 \\hmp$, and homogeneity does not exist, at least for what concerns galaxies. In this extreme case a new picture for the global metric\\cite{jampsl00} is then necessary.  \\end{itemize}  For some questions the fractal structure leads to a radically new perspective and this is hard to accept. But it is based on the best data and analyses available. It is neither a conjecture nor a model, it is a fact. The theoretical problem is that there is no dynamical theory to explain how such a fractal Universe could have arisen from the pretty smooth initial state we know existed in the big bang. However this is a  different question. The fact that something can be hard to explain theoretically has nothing to do with whether it is true or not. Facing a hard problem is far more interesting than hiding it under the rug by an inconsistent  procedure. For example some interesting attempts to understand why gravitational clustering generates scale-invariant structures have been recently proposed by de Vega et al\\cite{devega1,devega2,devega3}. Indeed this will be the key point to understand in the future, but first  we should agree on how these new 3d data should be analyzed. In addition, the eventual crossover to homogeneity has also to be found with our approach. If for example homogeneity would really be found say at $  \\sim100 \\hmp$, then clearly all our criticism to the previous methods and results  still holds fully. In summary the standard method cannot be used neither to disprove homogeneity, nor to prove it. One has simply to change methods."
    },
    "0002/astro-ph0002312_arXiv.txt": {
        "abstract": "A model is proposed for gamma-ray bursts based upon a neutrino burst of $\\sim 10^{52}$ ergs lasting a few seconds above a heated collapsing neutron star.  This type of thermal neutrino burst is suggested by relativistic hydrodynamic studies of the compression, heating, and collapse of close binary neutron stars as they approach their last stable orbit, but may arise from other sources as well.  We present a spherically symmetric hydrodynamic simulation of the formation and evolution of the pair plasma associated with such a neutrino burst. This pair plasma leads to the production of $\\sim 10^{51} - 10^{52}$ ergs in $\\gamma$-rays with spectral and temporal properties consistent with many observed gamma-ray bursts. ",
        "introduction": "Understanding the origin of gamma-ray bursts (GRBs) has been a perplexing problem since they were first  detected almost three decades ago \\citep{kso73}.  The fact that GRBs are distributed isotropically \\citep{mfw+92} suggests a cosmological origin. Furthermore, arcminute burst locations from BeppoSax have revealed that at least some $\\gamma$-ray bursts involve weak X-ray, optical, or radio transients, and are of cosmological origin \\citep{1997IAUC.6584}.  The Mg I absorption and [O II] emission lines along the line of sight from the GRB970508 optical transient, for example, indicate a redshift $Z \\ge 0.835$ \\citep{1997IAUC.6655}.  The implied distance means that this burst must have released of order $_>^\\sim 10^{51}$ ergs in $\\gamma$-rays on a time scale $\\sim$ seconds.  This energy requirement has been rendered even more demanding by other events such as GRB971214 \\citep{kfw+98} which appears to be centered on a galaxy at redshift 3.42.  This implies that the energy of a $4\\pi$ burst would have to be as much as $ 3 \\times 10^{53}$ ergs, comparable to the visible light output of $\\sim 10^9$ galaxies.  Based upon the accumulated evidence one can can now conclude that the following  four features probably characterize the source environment:  1) If the total burst energies are in the range of $10^{51}-10^{52}$ ergs, then a beaming factor of 10 to 100 is necessary;  2) The multiple peak temporal structure of most bursts probably requires either multiple colliding shocks \\citep{rm94,kps98} or a single shock impinging upon a clumpy interstellar medium \\citep{mr93b,dm99}; 3) The observed afterglows imply some surrounding material on a scale of light hours; and 4) the presence of [O II] emission lines suggests that the bursts occur in a young, metal-enriched stellar population.  Some proposed sources for the production of GRBs include accretion onto supermassive black holes, AGN's, relativistic stellar collisions, hypernovae, and binary neutron star coalescence.  Each of these possibilities, however, remain speculative until realistic models can be constructed for their evolution.  In this paper we construct a model for GRBs  produced by energetic neutrino emission from a heated neutron star.  Our specific model for the emission derives from the relativistic compression and heating of neutron stars near their last stable orbit, however any scenario by which energetic neutrino emission above a neutron star can endure for several seconds (e.g. tidal heating, MHD induced heating, accretion shocks, etc) might also power the gamma-ray burst paradigm described herein.  Our model is as follows.  A compressionally heated neutron star emits thermal neutrino pairs which, in turn, annihilate to produce a hot electron-positron pair plasma.  We model the expansion of the plasma with a spherically symmetric relativistic hydrodynamics computer program.  This simplification is justified at this stage of the calculations since the rotational velocity of the stars is about one third of the sound speed in the $e^+e^-$ pair plasma.  We then analyze and compare the contributions of photons from $e^+e^-$ pair annihilation as well as from an external synchrotron shock as the plasma plows into the interstellar medium.  We show that the characteristic features of GRBs, i.e. total energy, duration and gamma-ray spectrum, can be accounted for in the context of this model. ",
        "conclusions": "In this paper we have argued that  heated neutron stars (perhaps by compression of close neutron-star binaries) are viable candidates  for the production of large, high entropy per baryon, $e^+e^-$ pair plasma fireballs, and thus, for the creation of gamma-ray bursts.  We find that fireballs of total energy $E \\sim 10^{51}$ to $3 \\times 10^{52}$ ergs and an entropy per baryon of $s/k \\sim 10^5 - 10^6$ are possible.  Values for the  entropy as high as $10^7$ may be realized during the peak $\\nu \\overline{\\nu}$ luminosity (of $\\approx 10^{53}$ ergs sec$^{-1}$). Emergent gamma-rays yield a quasi-thermal spectrum peaked at $\\sim 100$ keV with an efficiency of conversion from pair plasma to photons of $\\sim 30$\\%. The lower entropy component of the fireball will initiate a shock which propagates into the ISM, generating an external shock GRB.  The calculation utilizing the supernova computer program (Section \\ref{neutrinoannihilation}) to describe the neutrino and matter transport, produces a baryonic wind that contains $\\sim 90$\\% neutrons. The decay of the neutrons to protons occurs on the same time scale as that for which the protons are decelerated by the intersteller medium. This delayed conversion of neutrons to protons will broaden the gamma-ray signal by a factor of a few. In addition, the decay electrons will strongly increase the entropy of the expanding plasma at the late times. In future work we  will quantify the role of neutrons and explore the possibility of fireball photons inverse-Compton up-scattering off of the accelerated electron distribution of the external shock. This corresponds to the emission scenario put forward by \\citet{lksc97}.  As of yet this model is spherically symmetric. Thus, it  can only generate bursts with a smooth light-curve structure.  However, we expect a large variety of GRB morphologies with varied time structure due to: 1) three dimensional resolution of the plasma flow; 2) plasma instabilities due to increased heating of the deposited plasma with time; and 3) variation in the ratio of star mass in the NSB, effecting the relative compression and heating rate of each star.  In future work we will numerically model, in three dimensions, the flow of the $e^+e^-$ pair plasma in the midst of the orbiting neutron stars.  We have written a three-dimensional general relativistic hydrodynamic code to study this three-dimensional behavior.  In particular, we wish to study the possible formation of jets along the orbital axis due to the collision of plasma blowing away from each star.  Also, we have done simulations which suggest that the internal magnetic field of the neutron stars may be high. Thus, the inclusion of magnetohydrodynamic plasma effects including Alfv\\'en instabilities and reconnections may ultimately be necessary.   The authors wish to thank the late Jean-Alain Marck for his inciteful and encouraging remarks on an earlier version of this manuscript. This work was performed under the auspices of the U.S. Department of Energy by University of California Lawrence Livermore National Laboratory under contract W-7405-ENG-48.  J.R.W. was partly supported by NSF grant PHY-9401636.  Work at University of Notre Dame supported in part by DOE grant DE-FG02-95ER40934, NSF grant PHY-97-22086, and by NASA CGRO grant NAG5-3818.  \\\\"
    },
    "0002/astro-ph0002349_arXiv.txt": {
        "abstract": "We review the current status of different mode-identification techniques based on observed line-profile variations. Three basic methods are currently available to identify the non-radial pulsation modes. These three methods are described, together with their different variants. We further present applications to real data, focusing especially on $\\delta\\,$Scuti stars. After having discussed all the methods and their applications, we present an inventory in which we compare the properties of the different methods. Finally, we end with future prospects, both on the theoretical and observational side. ",
        "introduction": "Intrinsic variable stars are an important diagnostic to test stellar models, which provide in turn valuable clues to our understanding of stellar and galactic evolution. Variable stars pulsate in so-called {\\it non-radial pulsation modes}.  Since the early seventies, ample observational evidence of the presence of such non-radial pulsations has become available.  From then on, detailed observational and theoretical studies of non-radial pulsations have been conducted. It has become clear that by understanding the pulsations in full detail, one can probe the internal structure of the stars and hence confront the results with stellar evolution theories, i.e., one can apply {\\it asteroseismology}.  The non-radial pulsations lead to periodic variations of physical quantities, such as surface brightness, radial velocity, temperature, etc.  By comparing the observed variations with those predicted by theory, it is in principle possible to determine the most important parameters that characterise the pulsation.  One specific aspect of studying pulsations therefore is to know what kind of modes are active in pulsating variables, i.e., the aspect of {\\it mode identification}.  Specifically, mode-identification techniques try to assign values to the spherical wavenumbers $(\\ell,m)$: the degree and azimuthal number of the spherical harmonic $Y_{\\ell}^m$ that describes the non-radial pulsation.  The identification of non-radial pulsation modes from observational data of variable stars is important, since it is a first step towards asteroseismology.  Indeed, the amount of astrophysical information that can be derived from the study of non-radial pulsations depends directly on the number of modes that can be successfully identified.  Inspired by this potential of  identifying the non-radial pulsation modes, the study of {\\it mode-identification techniques} has become an extended topic by itself in variable star research. We have been involved in the development and the application of such mode-identification methods.  In this review, we focus on the different identification techniques that are currently used to identify non-radial pulsations from line-profile variations.  The introduction of high-resolution spectrographs with sensitive detectors has had an enormous impact on the field of mode identification.  Spectroscopic data usually offer a very detailed picture of the pulsation velocity field. On the other hand, they require large telescopes and sophisticated instrumentation. Before accurate detectors were available, identifications had to be obtained from photometric observations.  These kind of data are suitable to study long-period pulsations because they can be obtained with small telescopes, which are available on longer time scales. For a review on photometric studies of $\\delta\\,$Scuti stars we refer to Poretti and to Garrido (these proceedings).  The plan of our paper is as follows. We first briefly give a (non-mathematical) introduction into the domain of non-radial pulsations. Next, we explain how a theoretical line profile can be calculated for a non-radially pulsating star.  The following section is devoted to the description of the different identification techniques: line-profile fitting, Doppler Imaging, and the moment method. We briefly review the main characteristics of each method in Section\\,5 and we finally end with some future prospects in this field of research. In particular, we point out the importance of the identification of pulsation modes in $\\gamma\\,$Dor stars. ",
        "conclusions": "In this review, we have discussed the different mode-identification techniques that are currently used to study the non-radial pulsations in pulsating stars from observations of line-profile variations. Three basic methods are presented, are compared to each other and applications to real observations of $\\delta\\,$Scuti stars are described.  Line-profile variations offer a very detailed picture of the various aspects of the pulsation velocity field. On the other hand, photometric observations are easier to obtain on a long time-basis and are as such often superior for a very accurate determination of the pulsation frequencies, especially in the case of lower-degree modes. High-degree modes hardly show up in photometry and can only be found from high-quality line-profile variation data. An example of additional modes being seen in spectra compared to photometry is presented by De Mey et al.\\ (1998). They have analysed high-quality line profiles of the multiperiodic $\\delta\\,$Scuti primary of the double-lined spectroscopic binary $\\theta\\,$Tuc. The dominant frequency is the same in the photometric and spectroscopic data, but the second frequency that shows up in the spectra was never found before in photometry. This example confirms that the gathering of simultaneous photometry and spectroscopy is the best strategy to find a complete and accurate identification of all the appearing modes in multiperiodic stars.  Future possible improvements from a theoretical point of view concern on the one hand the development of mathematical expressions for the phase and amplitude in Fourier space in such a way that these quantities can be immediately interpreted in terms of the physical parameters of the pulsational velocity field. We also briefly mention that a new method of ``Doppler Mapping'' was recently presented by Berdyugina et al.\\ (2000). They apply a spectral inversion technique to obtain maps of the surface corotating with the dominant pulsation mode. From these maps, they determine the pulsation degree and study the latitudinal distribution of the pulsation field.  The method still needs to be further explored. Secondly, the inclusion of temperature variations during the pulsation cycle is still not accurately done, since an adiabatic pulsation is assumed while it is to be expected that non-adiabatic effects are important in the outer region of the atmosphere where the observed spectral lines are formed. Finally, the inclusion of centrifugal forces may be an improvement for the most rapid rotators. The latter is only necessary for stars rotating close to their break-up velocity.  From an observational point of view, large progress can be expected in the near future now that better and better detectors become available. For the short-period $\\delta\\,$Scuti stars, the major problem in obtaining high temporal and spatial resolution profiles is that the ratio of the integration time to the main pulsation period is rather high. This was one of the reasons why the application of the moment method to the stars FG\\,Vir and X\\,Caeli was not very successful. The abovementioned ratio in these cases was respectively 13\\% and 8\\%, while it amounts to only 1\\% for our profiles of $\\rho\\,$Puppis shown in Figure\\,4.  For such high ratios, the pulsational motion is averaged out over a part of the cycle and this prevents unambiguous identifications, especially for multiperiodic stars.  An interesting new technique for the interpretation of line-profile variations is by working with cross-correlation functions instead of real spectra. Such a technique can be performed by means of current spectrographs such as ELODIE attached to the 1.93m telescope in the Haute-Provence Observatory. Our analysis of 20\\,CVn (Mathias \\& Aerts 1996) was already based on cross-correlation profiles and has shown that they perfectly contain the pulsational motion on the condition that the correlation is based on a suitable set of selected spectral lines. By using a cross-correlation function, one can significantly decrease the integration time and still obtain a high S/N ratio. At the same time, one can observe optically fainter stars with success.  More accurate versions of ELODIE-type spectrographs are CORALIE, attached to the Swiss 1.2m telescope and FEROS, attached to the ESO 1.5m telescope, both situated at La Silla in Chile.  Finally, we would like to point out that mode identification from line-profile variations will become an important tool to obtain some information on the nature of the excited modes in stars belonging to the new class of $\\gamma\\,$Dor stars. For reviews on this new group of pulsating stars we refer to Krisciunas (1998) and to Zerbi (these proceedings).  Since the multiperiodic variations detected in them have periods roughly a factor 20 longer than the period of the radial fundamental mode for such stars, high-order $g$-modes are believed to be the cause of the variability. However, there is yet no pulsation mechanism that can explain the onset and the maintenance of the pulsations in these stars.  Handler \\& Krisciunas have given subsequent updated lists of {\\it bona fide} members of the group, which currently constitutes 13 members.  We have recently taken the first steps towards the discovery of cool $g$-mode pulsators by searching for $\\gamma\\,$Dor stars in an unbiased sample of 39 new variable A2--F8 stars discovered by means of the Hipparcos mission (Aerts et al.\\ 1999b).  We have reported the discovery of 14 new $\\gamma\\,$Doradus variables among this unbiased sample.  We primarily focussed on the limited group of new variables for which both Hipparcos and Geneva data are available, mainly because the latter allow an accurate determination of the effective temperature. It is very likely, however, that our more extended list of 200 unclassified variable A2--F8 stars of which no Geneva data are at our disposal contains more objects of this type.  This seems to be confirmed by a recent analysis by Handler (1999).  In 1996, we also started a search for new $\\gamma\\,$Dor stars by means of ground-based Geneva photometry. Our search has resulted so far in the discovery of three new and some five suspected $\\gamma\\,$Dor stars (Eyer \\& Aerts, 2000).  In order to firmly establish the $\\gamma\\,$Dor nature of all these new candidates we have started a long-term spectroscopic campaign with CORALIE in the course of 1997, which is still ongoing. We found line-profile variability from our high-resolution spectra for almost all candidates. Some of them, however, turn out to be binaries (Eyer \\& Aerts, in preparation).  Line-profile studies of $\\gamma\\,$Dor stars are still scarce. Examples in which a large amount of spectra have been obtained and analysed are given by Balona et al.\\ (1996, $\\gamma\\,$Dor) and by Aerts \\& Krisciunas (1996, 9\\,Aur). The latter study is based on cross-correlations obtained with the (by now unmounted) spectrograph CORAVEL and showed convincingly that such correlation functions indeed contain a sufficient amount of information to characterise the pulsational behaviour.   It is clear that a combination of long-term photometry and spectroscopy is essential and the only way to study the multiperiodic variability in the $\\gamma\\,$Dor stars. The best observing strategy is the same as the one for the slowly pulsating B stars (Aerts et al.\\ 1999a), which are also multiperiodic $g$-mode pulsators. At present, we do not yet have a sufficient amount of line profiles for our targets, but we will continue our monitoring of $\\gamma\\,$Dor stars with CORALIE during the forthcoming years. This will eventually lead to mode identifications, by applying the moment method."
    },
    "0002/astro-ph0002380_arXiv.txt": {
        "abstract": "A survey of H$\\alpha$ emission in 320 spiral galaxies in 8 nearby clusters shows an enhancement of circumnuclear starburst emission with increasingly rich clusters.  These observations provide convincing evidence that spirals have been transformed into S0s in clusters predominantly by tidal forces, a picture fully in accord with the most recent numerical simulations of clusters.  For the richest clusters, the enhancement of starburst emission is greater than would be expected on the basis of increasing galaxy surface density alone, which may explain the anomalous result for the type--galaxy surface density (T--$\\Sigma$) relation found for low richness clusters at intermediate redshift. ",
        "introduction": "\\label{intro}  The remarkable changes in cluster disk galaxy populations between intermediate redshifts (z $\\sim$ 0.5) and the present are well known. Up to 50\\% of the population of intermediate redshift clusters are comprised of blue, star-forming galaxies, which have been shown to be predominantly normal spiral and irregular galaxies, a fraction of which are interacting or obviously disturbed.  By the present epoch, this population has been depleted by a factor of 2 and replaced by a population of S0 galaxies.  However the processes by which this has occurred are still not fully understood (cf. Dressler 1980; Dressler et al. 1997)  The processes which cause the transformation of the cluster spiral galaxy population to S0s are expected to have very significant effects on cluster galaxy star formation rates (SFRs).  We have undertaken a comparison of SFRs between field and cluster spirals in 8 low redshift clusters in order to investigate whether these SFRs show evidence of continuing morphological transformation of disk galaxies at the present epoch. The comparison survey, details of which are published elsewhere (Moss, Whittle \\& Irwin 1988; Moss \\& Whittle 1993; Moss, Whittle \\& Pesce 1998; Moss \\& Whittle 1999) uses H$\\alpha$ emission, resolved into disk and circumnuclear emission, as an estimator of the SFRs. ",
        "conclusions": "Although Lavery \\& Henry (1988) first proposed that the Butcher-Oemler effect could be explained as star formation triggered by galaxy--galaxy interactions in intermediate redshift clusters, it was long considered that the typical cluster velocity dispersion ($\\sim$ 1000 ${\\rm km}$ ${\\rm s}^{-1}$) was too high for strong tidal interactions to occur.  However recent work (e.g. Gnedin 1999) has shown that a non-static cluster potential can enhance tidal interactions for cluster galaxies.  Such a non-static potential can arise in sub-cluster merging, and indeed there is evidence that the two richest clusters in our sample (Abell 1367 and 1656) are recent post-merger systems (Donnelly et al. 1998; Honda et al. 1996).  It appears that the merger events in these clusters leading to a rapidly varying cluster potential, may have caused an increase in galaxy tidal interactions and the associated observed enhancement of circumnuclear starburst emission.  Tidal interactions of galaxies in clusters are likely to be an effective mechanism for the transformation of spirals to S0 galaxies (e.g. Gnedin 1999).  It was noted above that the observed enhancement of circumnuclear emission in spirals with increasing cluster density is not wholly accounted for by that due to increasing local galaxy surface density.  This implies that tidal interactions, and associated morphological transformation of spirals to S0s, proceeds faster in richer as compared to poorer clusters, perhaps because sub-cluster merging is more common for the former. This in turn may explain the hitherto anomalous absence of a type--local galaxy surface density ($T-\\Sigma$) relation in irregular clusters at intermediate redshift (Dressler 1980; Dressler et al. 1997). Whereas significant morphological transformation of the cluster disk galaxy population may be expected for regular (rich) clusters at $z \\sim 0.5$ (for which the timescale for transformation is shorter) and for irregular (poor) clusters at $z \\sim 0$ (for which a longer time duration for transformation is available),  for irregular clusters at $z \\sim 0.5$, there has been insufficient time for this to take place, leading to the observed absence of a $T-\\Sigma$ relation."
    },
    "0002/astro-ph0002455_arXiv.txt": {
        "abstract": "We calculated models of massive AGB stars with a self-consistent coupling of time-dependent mixing and nuclear burning for 30 isotopes and 74 reactions. Overshoot with an exponentially declining velocity field was considered and applied during all stages of evolution and in all convective regions. Very efficient $3^{\\rm rd}$ dredge-up was found even overcompensating the growth of the hydrogen-exhausted core after a few thermal pulses. Hot bottom burning occurs for $M\\ge 4\\,$M$_{\\odot}$ within the sequences with overshoot. Carbon star formation in these more massive AGB stars is delayed or even prevented by hot bottom burning despite the very efficient dredge-up. With the simultaneous treatment of mixing and burning the formation of Li-rich AGB stars due to the Cameron-Fowler mechanism was followed. For  a $6\\,$M$_{\\odot}$ model the maximum Li abundance was found to be $\\epsilon(^7{\\rm Li}) \\approx 4.4$. ",
        "introduction": "In recent stellar evolution calculations Herwig et al.\\ (1997) have considered diffusive overshoot for all convective boundaries which leads to a considerable change in the models. This method provides for Asymptotic Giant Branch (AGB) stars a sufficient amount of dredge-up to form low-mass carbon stars as required from observations. Additionally it leads to the formation of $^{13}$C as a neutron source to drive the s-process in these stars. More details are given in Herwig et al.\\ (1999).  However the consequences for massive AGB stars, which suffer from hot bottom burning, remain to be investigated in detail. Hot bottom burning, i.e.\\ the penetration of the convective envelope into the hydrogen burning shell, provides lithium-rich AGB stars and may delay or prevent the  carbon star stage by turning $^{12}$C into $^{13}$C and  $^{14}$N. We have extended our computations of AGB stars with diffusive overshooting to hot bottom burning models including a self-consistent coupling of nucleosynthesis and mixing. ",
        "conclusions": "We calculated models of massive AGB stars (solar metallicity) with a self-consistent coupling of time-dependent mixing and nuclear burning for 30 isotopes and 74 reactions. Treating exponential overshoot (Freytag et al. 1996) as in Herwig et al.\\ (1997)  and applying it to all convective regions we find very efficient $3^{\\rm rd}$ dredge-up with $\\lambda > 1$ after a few thermal pulses. Overshoot leads to intershell dredge-up and provides considerably changes of the intershell abundances strongly enriched in carbon and oxygen, i.e. [He,C,O]=[40,40,14] by mass. These  intershell abundances determine the strength of the third dredge up, whereas envelope overshoot is mainly required to penetrate the H/He discontinuity. Temperatures at the bottom of the flash-driven convection zone of the helium burning shell are in excess of $300 \\cdot 10^{6}$\\,K after a few pulses activating the $^{22}$Ne($\\alpha$,n)$^{25}$Mg reaction for $s$-process nucleosynhesis.  Hot bottom burning occurs for $M\\ge 4\\,$M$_{\\odot}$ within our sequences with overshoot. Only the  $4\\,$M$_{\\odot}$ model becomes a carbon star whereas carbon star formation in more massive AGB stars is delayed or even prevented by hot bottom burning despite of the very efficient dredge-up. With the simultaneous treatment of mixing and burning the formation of Li-rich AGB stars due to the Cameron-Fowler mechanism was studied. Already the  $4\\,$M$_{\\odot}$ model becomes lithium-rich leading to the formation of a lithium-rich carbon star. For  $M\\ge 6\\,$M$_{\\odot}$ the maximum Li abundance was found to be $\\epsilon(^7{\\rm Li}) \\approx 4.4$. \\vspace*{-1ex}"
    },
    "0002/astro-ph0002039_arXiv.txt": {
        "abstract": "The white dwarf Giclas 29-38 has attracted much attention due to its large infrared excess and the suggestion that excess might be due to a companion brown dwarf.  We observed this object using speckle interferometry at the Keck telescope, obtaining diffraction-limited resolution (55 milliarcseconds) at K band, and found it unresolved.  Assuming the entire K band excess is due to a single point-like companion, we place an upper limit on the binary separation of $30$ milliarcseconds, or $0.42$ AU at the star's distance of 14.1 pc.  This result, combined with astroseismological data and other images of G29-38, supports the hypothesis that the source of the near-infrared excess is not a cool companion but a dust cloud. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002513_arXiv.txt": {
        "abstract": "Observations of the damped \\lya systems provide direct measurements on the chemical enrichment history of neutral gas in the early universe.  In this Letter, we present new measurements for four damped \\lya systems at high redshift.  Combining these data with $\\lbrack$Fe/H$\\rbrack$ values culled from the literature,  we investigate the metallicity evolution of the universe from $z \\approx 1.5-4.5$. Contrary to our expectations and the predictions of essentially every chemical evolution model, the $\\N{HI}$-weighted mean $\\lbrack$Fe/H$\\rbrack$ metallicity exhibits minimal evolution over this epoch. For the individual systems, we report tentative evidence for an evolution in the unweighted $\\lbrack$Fe/H$\\rbrack$ mean and the scatter in $\\lbrack$Fe/H$\\rbrack$ with the higher redshift systems showing lower scatter and lower typical $\\lbrack$Fe/H$\\rbrack$ values. We also note that no damped \\lya system has $\\lbrack$Fe/H$\\rbrack$ $< -2.7$~dex. Finally, we discuss the potential impact of small number statistics and dust on our conclusions and consider the implications of these results on chemical evolution in the early universe. ",
        "introduction": "The damped \\lya systems, neutral hydrogen gas layers identified in the absorption line spectra of background quasars, dominate the neutral hydrogen content of the Universe at all epochs.   At high redshift, these systems are widely accepted as the progenitors of present-day galaxies for the following reasons: (i) their very large HI column densities, $\\N{HI} > N_{thresh} = 2 \\times 10^{20} \\cm{-2}$, imply overdensities $\\delta \\rho / \\rho \\gg 100$, i.e., these are virialized systems at high redshift; (ii)  they contain the majority of neutral gas in the early Universe and are therefore the reservoirs for galaxy formation; (iii) their gas density $\\Omega_{gas}$ at redshift $z \\approx 2 -3$ is consistent with the mass density of stars today (\\citep{wol95}). While the physical nature of the damped \\lya systems is still controversial (\\citep{pro97,hae98,maller99,lebr97}), by studying the chemical abundances of the damped \\lya system one directly traces the chemical enrichment history of the Universe at high redshift. Observing damped \\lya systems is equivalent to poking sightlines through the ISM of protogalaxies. Because these observations are biased to HI cross-section and the HI gas mass of a system is proportional to $\\int \\sigma \\cdot N$, one can measure global properties of the universe simply by weighting the measurement from each damped system by $\\N{HI}$. At the same time,  the observations afford an efficient means for examining the characteristics of individual protogalaxies in the early Universe. In this Letter, we examine the metallicity of the damped \\lya systems from $z \\approx 2 - 4.5$ which places tight constraints on chemical evolution models (e.g.\\ Pei et al.\\ 1999), as well as a valuable consistency check on SFR observations (\\citep{ptt99a}).  Over the past decade, several groups have surveyed the metallicity of the damped \\lya systems from $z \\approx 1 - 4$ (\\citep{ptt94,ptt97,ptt99a,lu96,pro99}). To date, the chemical abundances of over 40 systems have been measured, the majority with $z = 1.5 - 3$ where the identification and follow-up observations of damped \\lya systems is most efficient. These studies argue that at $z \\approx 2$, the mean metallicity of the damped systems is approximately 1/10 $-$ 1/30 solar metallicity ([Zn/H]~$\\approx -1.1$, [Fe/H] $\\approx -1.5$) with a large scatter from nearly solar to less than $1/100$ solar metallicity.  At very high redshift ($z>3$), the picture is far less certain.  Focusing on a sample of seven $z>3$ damped \\lya systems, Lu et al.\\ (1996,1997) noted that the metallicity of these systems is significantly lower than the $z<3$ observations. In turn, the authors argued that $z \\approx 3$ marked the epoch where significant star formation begins, a claim with important implications for the processes of galaxy formation.  In this Letter we present new measurements on the metallicity of four damped \\lya systems (including three at $z>3.5$) and together with the data from Prochaska \\& Wolfe (1999) double the sample of $z>3$ systems. The new full sample -- including the systems from Lu et al.\\ (1996,1997) -- reveals evidence for little change in the $\\N{HI}$-weighted mean metallicity of the neutral Universe from $z \\approx 1.5-4.5$, contrary to the predictions of essentially every chemical evolution model.   On the other hand, we find tentative evidence for an evolution in the unweighted mean and scatter of [Fe/H] for individual damped \\lya systems. Finally, we comment on the robustness of these results (particularly in the light of small number statistics and dust), speculate on the implications for chemical enrichment, and discuss the prospects for future advances. ",
        "conclusions": " (i) star formation proceeds rapidly ($< 10^7$ years) to bring the metallicity to 1/100 solar after the formation of a damped system; (ii) either the damped system or its progenitors have been undergoing star formation for a lengthy time; and/or (iii) a generation of Population III stars has pre-enriched all of the gas.  The future prospects for improving the $z>3$ observational sample are excellent.  While further progress with HIRES and similar instruments is limited by the faintness of $z>4$ quasars, the new ESI instrument on Keck~II will be ideal for surveying the $\\N{HI}$ and metal content of a large ($N>20$) sample of very high redshift damped \\lya systems. We intend to pursue such a program over the next few years, taking full advantage of the ever increasing sample of known $z>4$ quasars (\\citep{fan99a})."
    },
    "0002/astro-ph0002275_arXiv.txt": {
        "abstract": "In order to ascertain the extent of metal enrichment in the Ly$\\alpha$ forest, we have analysed a very high S/N spectrum of the $z = 3.625$ QSO Q1422+231.  We find that in high column density Ly$\\alpha$ clouds, the power law column density distribution function of C~IV continues down to log $N$(C~IV) = 11.7.  In addition, by analysing pixel-by-pixel optical depths we show that there are considerably more metals in the Ly$\\alpha$ forest than are currently directly detectable. ",
        "introduction": "Our understanding of the Ly$\\alpha$ forest and its connection with the Intergalactic Medium (IGM) have undergone radical revision in recent years.  The paradigm in which the IGM consists of discrete, isolated clouds has been revolutionised by a generation of hydrodynamical simulations (e.g. Hernquist et al 1996).  These simulations show that the `bottom-up' hierarchy of structure formation knits a complex but smoothly fluctuating cosmic web, consisting of filaments, knots and extensive `voids'.  Originally thought to be chemically pristine, it is now well-established that a large fraction of the strongest Ly$\\alpha$ absorbers exhibit some metal enrichment, most notably C~IV (e.g. Cowie et al 1995). Here, we address two specific questions.  Firstly, are the C~IV absorbers that have thus far been detected just the tip of the iceberg and can more sensitive spectra mine ever weaker systems?  Secondly, to what H~I column densities does the enrichment extend?  This latter point has particularly poignant implications for the origin and transport mechanism of these metals. In-situ formation in a nearby galaxy could be responsible for the enrichment of its local IGM, explaining the presence of C~IV in the relatively high column density Ly$\\alpha$ clouds.  However, low column density Ly$\\alpha$ clouds (log$N$(H~I) $<$ 14.0), which are associated with physically less dense regions, are found further from the sites of star formation.  The presence of C~IV in these clouds could be indicative of widespread metal enrichment, possibly by an early epoch of Population III star formation.  This is an overview of work described in more detail in Ellison et al (2000). ",
        "conclusions": ""
    },
    "0002/astro-ph0002043_arXiv.txt": {
        "abstract": "We observed Sirius, Altair, and Procyon with the NICMOS Coronagraph on the Hubble Space Telescope to look for scattered light from exo-zodiacal dust and faint companions within 10 AU from these stars.  We did not achieve enough dynamic range to surpass the upper limits set by IRAS on the amount of exo-zodiacal dust in these systems, but we did set strong upper limits on the presence of nearby late-type and sub-stellar companions. ",
        "introduction": "Several main sequence stars are close enough that a large telescope operating at the diffraction limit can resolve the terrestrial planet-forming region within 10 AU from the star \\markcite{kuch98} (Kuchner, Brown \\& Koresko 1998).  We used the NICMOS coronagraph to image three of the nearest main-sequence stars in the near-infrared to look for circumstellar material---exozodiacal dust and faint companions---in this relatively uncharted circumstellar region.  Any dust orbiting close to one of our targets must have been generated recently by some population of larger bodies, since dust near a star quickly spirals into the star due to Poynting-Robertson drag (Robertson 1937).  Ten micron diameter dust 3 AU from a G star spirals into the star on time scales of $\\sim10^5$ years; for brighter stars, this timescale is shorter.  Around the sun, zodiacal dust forms when asteroids collide and when comets outgass; a search for exozodiacal dust is therefore implicitly a search for extra-solar asteroid or comet-like bodies that make dust.  Several disks around nearby main-sequence stars appear to have exozodiacal components.  Some, such as the disk associated with Beta-Pictoris, reveal their warm dust as a silicate emission feature at 10 microns \\markcite{tele91} (Telesco \\& Knacke 1991).  Others, like the disk around HR 4796, show resolved emission at 10 microns that is interpreted as exozodiacal \\markcite{koer98} (Koerner et al. 1998).  Dust clouds like these, which have $\\sim 1000$ times as much warm dust as our sun, emit thermal radiation substantially in excess of the stellar photospheric emission, and can often be detected photometrically by studying the spectral energy distribution of the star in the mid-infrared.  However, many less massive exozodiacal clouds may never be detectable photometrically because no stellar spectrum is known to better than $\\sim$ 3\\% in the mid-infrared \\markcite{cohe96}(Cohen et al. 1996).  We have begun to search for disk that are too faint to be detected photometrically by spatially resolving the critical regions less than 10 AU from nearby stars.  Coronagraphic images can also reveal faint companions to nearby stars.  Such companions can go undetected by radial velocity surveys because of their small masses or long orbital periods. Siruis and Procyon both have white dwarf companions whose orbits are well studied, but analyses of the orbital motion in these systems leave room for additional low mass companions.  The Sirius system in particular, so prominent in the night sky, has spurred much debate in the last century over its properties.  Three analyses of the proper motion of Sirius have suggested that there may be a perturbation in the orbit of Sirius B with a $\\sim 6$ year period \\markcite{vole32}\\markcite{walb83}\\markcite{bene95} (Volet 1932; Walbaum \\& Duvent 1983; Benest \\& Duvent 1995). These analyses do not indicate whether the perturbing body orbits Sirius A or B, and dynamical simulations indicate that stable orbits exist around both Sirius A and B at circumstellar distances up to more than half the binary's periastron separation \\markcite{bene89}(Benest 1989).  If such a companion were in a simple face-on circular orbit it would appear at a separation of 4.2 AU (1.6 arcsec) from Sirius A or a separation of 3.3 AU (1.3 arcsec) from Sirius B assuming that the masses for Sirius A and B a are 2.1 and 1.04 M${}_{\\odot}$ respectively \\markcite{gate78}(Gatewood \\& Gatewood 1978).  Benest \\& Duvent \\markcite{bene95}(1995) do not derive a mass for the hypothetical companion from observations of the system, but they estimate that a perturber much more massive than 0.05 M${}_{\\odot}$ would rapidly destroy the binary.  Perhaps the most interesting debate about Sirius is whether or not the system appeared red to ancient observers $\\sim 2000$ years ago. Babylonian, Graeco-Roman and Chinese texts from this time period have separately been interpreted to say that Sirius was a red star \\markcite{brec79}\\markcite{schl85}\\markcite{bonn91} (Brecher 1979; Schlosser \\& Bergman 1985, Bonnet-Bidaud \\& Gry 1991). \\markcite{tang86}\\markcite{vang84}\\markcite{mccl87} Tang (1986), van Gent (1984) and McCluskey (1987) have attacked some of these reports, claiming that they represent mistranslations or misidentifications of the star.  However, if the Sirius system did indeed appear red, the existence of a third star in the group interacting periodically with Sirius A could explain the effect \\markcite{bonn91}(Bonnet-Bidaud 1991).  The low mass companions ($< 0.1$ M${}_{\\odot}$) that we could hope to detect with NICMOS are late-type stars or warm brown dwarfs, shining with their own thermal power in the near-infrared.  Schroeder and Golimowski \\markcite{schr96}(1996) recently imaged Sirius, Procyon and Altair at visible wavelengths with the Planetary Camera on HST in a search for faint companions.  Our observations  are more sensitive to late-type companions because of the high dynamic range of the NICMOS coronagraph, and because these objects are brighter in the infrared than the optical. ",
        "conclusions": ""
    },
    "0002/astro-ph0002333_arXiv.txt": {
        "abstract": "In the course of a Cycle~8 snapshot imaging survey with STIS, we have discovered that the z=1.565 quasar HE~0512$-$3329 is a double with image separation 0\\farcs644, differing in brightness by only 0.4~magnitudes.  This system is almost certainly gravitationally lensed.  Although separate spectra for the two images have not yet been obtained, the possibility that either component is a Galactic star is ruled out by a high signal-to-noise composite ground-based spectrum and separate photometry for the two components: the spectrum shows no trace of any zero redshift stellar absorption features belonging to a star with the temperature indicated by the broad band photometry.  The optical spectrum shows strong absorption features of \\MgII, \\MgI, \\FeII, \\FeI, and \\CaI, all at an identical intervening redshift of z=0.9313, probably due to the lensing object.  The strength of \\MgII\\ and the presence of the other low-ionization absorption features is strong evidence for a damped \\Lya\\ system, likely the disk of a spiral galaxy.  Point spread function fitting to remove the two quasar components from the STIS image leads to a tentative detection of a third object which may be the nucleus of the lensing galaxy.  The brighter component is significantly redder than the fainter, due to either differential extinction or microlensing. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002105_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "Remarks on the Adopted Indexing Scheme}  One of the advantages of using separated index and list files is that the size of the files that are accessed most frequently by the search engine (the list of bibcodes and the index files) is kept small so that their contents can be loaded in random access memory and searched efficiently (SEARCH). For instance, the size of the text index file for the astronomy database is approximately 16 MB, and once the numerical entries are converted into binary representation when loaded in memory by the search engine, the actual amount of memory used is less than 10 MB.  The use of integer sequential identifiers in the list files allows more compact storage of the document identifiers as well as implementation of fast algorithms for merging search results (since all the operations are executed in 32-bit integer arithmetic rather than having to operate on 19-character strings). For instance, recent indexing of the ADS astronomy database produces text inverted files which have sizes approaching 500 MB, while the size of the text list file is about 140 MB.  The choice of a word weight which is a function of only the document frequency allows us to store word weights as part of the index files.  It has been shown that a better measure for the relevance of a document with respect to a query word is obtained by taking into account both the document frequency $df$ and the term frequency $tf$, defined as the frequency of the word in each document in which it appears (\\cite{SALTON88}), normalized to the total number of words in the document.  The reasoning behind this is that a word occurring with high relative frequency in a document and not as frequently in the rest of the database is a good discriminant element for that document. Although we had originally envisioned incorporating document-specific weights in the list files to take into account the relative term frequency of each word, we found that little improvement was gained in document ranking. This is probably due to the fact that the collection of documents in our databases is rather homogeneous as far as document length and characteristics are concerned.  Eventually the choice was made to adopt the simpler weighting scheme described above.  The procedures used to create the inverted files can scale well with the size of the database since the global inverted file is always created by joining together partial inverted files. This allows us to limit the number of hash entries used by the indexer program during the computation of the inverted files. According to Heap's law (\\cite{HEA78}), and as verified experimentally in our databases, a body of $n$ words typically generates a vocabulary of size $V = K n^{\\beta}$ where $K$ is a constant and $\\beta \\approx 0.4 - 0.6$ for English text (\\cite{NAV98}). Since the size of the vocabulary $V$ corresponds to the number of entries in a global hash table used by the indexing software, we see that an ever-increasing amount of hardware resources would be necessary to hold the vocabulary in memory; our choice of a partial indexing scheme avoids this problem. Furthermore, the incremental indexing model is quite suitable to being used in a distributed computing environment where different processors can be used in parallel to generate the partial inverted files, as has been recently shown by \\cite{KITA97}.  The procedures used to create the list and index files make use of memory sparingly, so that processing of entries from the occurrence tables is essentially sequential. The only exception to this is the handling of groups of synonyms. In that case, the data structures used to maintain the entries for the words in the current synonym group are kept in memory while the cumulative list of sequential identifiers for the entire group is built.  The memory is released as soon as the entries for the current synonym group are written to the list and index files."
    },
    "0002/astro-ph0002111_arXiv.txt": {
        "abstract": "We have undertaken an imaging survey of 34 nearby galaxies in far--ultraviolet (FUV, $\\sim 1500 \\AA$) and optical ($UBVRI$) passbands to characterize galaxy morphology as a function of wavelength. This sample, which includes a range of classical Hubble types from elliptical to irregular with emphasis on spirals at low inclination angle, provides  a valuable database for comparison with images of high--$z$ galaxies whose FUV light is redshifted into the optical and near--infrared bands. Ultraviolet data are from the UIT {\\it Astro--2} mission. We present images and surface brightness profiles for each galaxy, and we discuss the wavelength--dependence of morphology for different Hubble types in the context of understanding high--$z$ objects. In general, the dominance of young stars in the FUV produces the patchy appearance of a morphological type later than that inferred from optical images. Prominent rings and circumnuclear star formation regions are clearly evident in FUV images of spirals, while bulges, bars, and old, red stellar disks are faint to invisible at these short wavelengths. However, the magnitude of the change in apparent morphology ranges from dramatic in early--type spirals with prominent optical bulges to slight in late--type spirals and irregulars, in which young stars dominate both the UV and optical emission. Starburst galaxies with centrally concentrated, symmetric bursts display an apparent ``E/S0'' structure in the FUV, while starbursts associated with rings or mergers produce a peculiar morphology. We briefly discuss the inadequacy of the optically--defined Hubble sequence to describe FUV galaxy images and estimate morphological $k$--corrections, and we suggest some directions for future research with this dataset. ",
        "introduction": "\\par An understanding of the ultraviolet (UV) properties of local galaxies is essential for interpreting images of high--redshift systems and especially critical in searches for morphological evolution in the galaxy population. At redshifts of $ z \\sim 3 - 10$, the rest--frame UV light of galaxies is shifted into the optical (1500$\\AA$ observed in $V$--band at $z \\sim 2.7$) and near--infrared (1500$\\AA$ observed in $H$--band at $z \\sim 10$) regions of the spectrum that are easily observable from the ground.  Although one operational approach to overcoming this ``band--shifting'' problem is to observe in the infrared ({\\it e.g.} \\cite{tep98}), the rate of data acquisition for galaxy images is generally highest at optical wavelengths where detectors are larger and the sky background is lower than in the infrared.  Under these circumstances, the most direct method to compare present--day galaxies to those in the early universe is to build a large sample of local UV galaxy data.  \\par Recent deep surveys suggest that galaxies with peculiar morphology are more prevalent at high redshift ({\\it e.g.} \\cite{bri98}; \\cite{abr96a}). However, it is difficult to quantify the influence of band--shifting on observed differences between the optical characteristics of nearby galaxies and the rest--frame UV appearance of distant ones because UV and optical morphology are not often well--coupled (\\cite{oco97a}; Marcum et al. 1997, 2000).  Restframe UV images of galaxies often suggest a later Hubble type than the corresponding optical appearance:  bulges are less prominent, galaxy light appears more patchy, and in extreme cases the different star--forming regions in a single galaxy may appear as separate systems in the UV (\\cite{oco97a}; \\cite{oco96}). On the other hand, recent deep NICMOS imaging of distant galaxies at longer rest wavelengths suggests that some of the morphological peculiarities are still visible in the optical and thus are intrinsic in the galaxies' structure ({\\it e.g.} \\cite{bun99}; \\cite{tep98}). Although methods exist to estimate UV morphology by extrapolation from optical data (\\cite{abr98}; \\cite{afm97}), they rely on the use of global galaxy spectral energy distributions  that may not accurately represent localized conditions, especially in dusty or starbursting regions. For example, Donas, Milliard, \\& Laget (1995) find a variation among UV$-b$ colors in galaxies with identical $b-r$. The availability of FUV images for bright, well--resolved galaxies of a range of Hubble types would facilitate a more direct comparison between local and distant galaxy populations.  \\par The appearance of galaxies in the UV is determined primarily by emission from hot stars and by the distribution of dust, with a contribution in the central regions from an AGN if one is present and unobscured. Light in the $\\sim$1500$\\AA$ far--ultraviolet (FUV) band of the Ultraviolet Imaging Telescope (UIT, \\cite{ste97}) originates in young objects of spectral types O and B (Fanelli et al. 1997) and, if old populations are present, in certain types of low--mass stars in late stages of evolution (the ``UVX'' population, see \\cite{oco99} for a review). Thus this region of the spectrum can probe timescales of young stellar populations between that studied via $H\\alpha$ emission (only O stars, $\\sim$ 5 Myr) and that evident in the optical colors (which trace the age of a stellar population on timescales of a few Gyr; O'Connell 1997b) .  \\par Opacity due to dust in the FUV is a factor of $\\sim$2.5 higher than in the optical $V$ band, assuming a Galactic extinction law (\\cite{ccm89}), and even larger for Small Magellanic Cloud--type dust (\\cite{gor97}).  The young, massive stars that provide the bulk of the UV emission also tend to have small scale heights, less than or similar to that of the dust, while older stars that dominate at optical wavelengths and have scale heights greater than that of the dust layer (\\cite{mih81}). The degree to which dust extinguishes the UV light of highly inclined disk galaxies is unclear. It likely depends strongly on the relative geometry of star--forming regions and dust concentrations, a phenomenon recently noted in optical/NIR studies of nearby galaxies as well as in theoretical studies of systems with a mixture of stars and dust (\\cite{kuc98}; Gordon et al. 1997; \\cite{hui94}).  For instance, NGC 4631 is quite bright in the FUV and blue in FUV--optical colors, and it does not show evidence for strong attenuation (Smith et al. 1997, 2000; \\cite{oco97a}). However, NGC 891 is only barely detected (Marcum et al 1997, 2000), a phenomenon that Smith et al. suggest is due to extinction effects.  UV morphology may also be dramatically influenced by specific features such as dust lanes, as noted in Cen~A by O'Connell (1997a). Scattering by dust grains is quite efficient in the FUV (\\cite{wit92}) and may contribute significantly to the emission near H II regions and in spiral arms (\\cite{wal97a}; \\cite{ste82}). \\par Prior to the Space Shuttle--borne UIT {\\it Astro} missions, UV imaging of galaxies was limited to low resolution and/or small fields of view (see \\cite{bro99} for a detailed review).  UIT--prototype rocket--borne experiments in the 1980's provided images of five nearby galaxies with a  resolution of $\\sim$15$\\arcsec$  and ten more with a resolution of 100$\\arcsec$ (\\cite{hil84}; \\cite{boh91}). The FOCA balloon experiment yielded images of similar resolution for six nearby galaxies (\\cite{ble90}).  Recently, higher resolution UV images of 110 galaxy nuclei (\\cite{mao96}) and nine starburst galaxies (\\cite{meu95}) have been obtained at 2200$\\AA$ using the Hubble Space Telescope's Faint Object Camera with a very small (22\\arcsec) field of view. A FUV ($\\sim$1500$\\AA$) survey of galaxies at redshifts of $\\sim$~0.1--0.2 is also underway using the Space Telescope Imaging Spectrograph, again with a small (25\\arcsec) field of view (\\cite{gar97}). With a 40$\\arcmin$ field of view and $\\sim$3$\\arcsec$ resolution, UIT was optimally suited to study the morphology of nearby galaxies with good resolution. The usefulness of UV and optical images and radial profiles to study morphology variations with wavelength was demonstrated  by \\cite{cor94} with UIT {\\it Astro--1} data for M74,  by \\cite{hil97} for M51, and  by \\cite{rei94} for M81.   \\par In this and a companion paper (Marcum et al. 2000), we present the results of an UV--optical imaging survey of nearby bright galaxies.  The UV images were obtained by the UIT during the {\\it Astro--1} and {\\it Astro--2} Spacelab missions in 1990 and 1995 respectively.  Ground--based optical CCD images of the sample galaxies were obtained with pixel scales and fields of view approximately comparable to those of the UIT images.  The {\\it Astro--1} data presented in Marcum et al. include both FUV ($\\sim$1500$\\AA$) and NUV($\\sim$2500$\\AA$) data as well as optical images for 43 galaxies.  This paper includes FUV and optical data for 34 galaxies observed during {\\it Astro--2} with Hubble types specifically chosen to span the range E to Irr, including some interacting/merging systems. There is some emphasis in the sample selection on face--on spirals in which spiral arm, bar, and ring morphologies are easily observed. In Section 2 we discuss the sample selection, describe the UIT and optical observations, and explain the basic data reduction. Section 3 contains details of the registration of FUV and optical images and the extraction of surface brightness and color profiles.  A qualitative comparison of UV and optical morphology, based on the images and profiles, is presented in Section 4.  In Section 5 we provide a brief summary and discuss avenues of future investigation, including studies of the morphological $k$--correction and analysis of the star--formation histories of individual galaxies.  In a subsequent paper (\\cite{kuc00}), we will investigate the quantitative concentration and asymmetry indices for these galaxies and consider the implications of these results for the study of high--redshift galaxies. ",
        "conclusions": "\\par We have presented FUV and optical images and surface brightness profiles for 34 galaxies observed with the UIT and compared the apparent morphology at different wavelengths in the context of interpreting images of high--redshift objects. The UIT data represent the most detailed set of FUV images obtained to date for large, nearby galaxies that are well--resolved. Present--day elliptical and S0 galaxies appear smooth and symmetric in FUV as well as in optical images, but they are much more compact in the FUV and may be faint and/or unresolved at high redshift unless they are actively forming stars at that particular look--back time. The majority of spiral galaxies in our sample appear to have a later Hubble type in the FUV than at optical wavelengths due to increased patchiness in combination with the fading of light from the bulge and/or bar populations. This effect is particularly dramatic for the early--type spirals, in which bulges and often bars are prominent optical features.  The optical and FUV morphologies of late--type spirals and irregular galaxies do not differ as much as the earlier types because young stars dominate the light in both spectral regimes.  Some galaxies appear highly fragmented in the FUV images in spite of a regular optical morphology, making it hard to determine from UV data alone whether they are single systems or multiple mergers. Even assuming their current rates and patterns of star formation, many spirals viewed at high redshift would be assigned later Hubble types than their local counterparts due to the bandshifting effects. Central or circumnuclear starbursts and star--forming rings, and bright spiral arms dominate the FUV light and FUV--optical color profiles of many galaxies.  Depending on the relative geometry of dust and young stars, extinction can play a varying role (small to dominant) in determining the apparent FUV morphology.  From the galaxy sample presented here, it is apparent that understanding the correlation between FUV and optical morphology depends on one's knowledge of not only the global optical Hubble type but also of smaller scale features, mostly those due to recent star--formation patterns.  \\par The fragmented appearance of many of the FUV images and the prominence of young stars in structures such as rings or fragments of spiral arms produces FUV morphologies that do not readily fit into the traditional bins of the Hubble sequence.  In order to characterize rest--frame UV galaxy populations at both low and high redshift, it will be necessary to find new ways of describing the morphology that do not by default cause the majority of galaxies to fall into the ``peculiar'' or other catch--all bins.  It would also be desirable to have an objective classification scheme to express the similarities or differences between local and distant galaxies in quantitative terms.  However, UV data pose problems for many automated classification schemes because galaxy centers are often ill--defined and the patchy nature complicates the selection of an aperture that encloses a single system.  Future work in this area will use UIT data to identify new features or indices that better describe the UV morphology of galaxies.  \\par The dataset presented in this paper will be valuable for studying the morphological $k$--correction that must be taken into account when interpreting deep optical images such as the {\\it Hubble Deep Field} (HDF, \\cite{wil96}).  The dependence of FUV morphology on features that are considered minor or secondary details in the optical highlights the difficulty of studying evolution in the galaxy population by comparing rest--frame UV observations of high--redshift galaxies to the optical properties of local samples.  The comparison between local and distant UV galaxy images is much more direct than relying on estimates of the UV morphology extrapolated from optical light and colors ({\\it e.g.} Abraham et al. 1997). Indeed, a small sample of UIT galaxy images from the {\\it Astro--1} mission have been used to simulate the appearance of high--redshift counterparts of nearby bright galaxies (Giavalisco et al. 1996b). In a subsequent paper (\\cite{kuc00}), we simulate the cosmological distance effects of surface brightness dimming and loss of spatial resolution on the FUV images of our larger {\\it Astro--2} galaxy sample. These simulations provide examples of how local galaxies would appear in the HDF at redshifts of $\\sim 1 - 4$.  In that paper, we also attempt to quantify the effects of bandshifting and distance on morphology using simple parameters such as the central concentration and asymmetry of galaxies.  \\par One of the primary goals of studying morphology is to infer the past or ongoing physical processes that have shaped galaxies. In combination with current knowledge about the dynamical states of nearby galaxies, our multi--wavelength dataset and future UV imaging data from the GALEX sky survey (\\cite{mtin97}) could be used to relate UV morphology to physical structures or properties.  For example, Waller et al. (1997b, 2000) find that UV--bright starburst rings likely result from dynamical resonances with a bar component. Thus ring features may predict the presence of an underlying disk and bar system even in UV images that appear to be fragmented or otherwise dynamically disorganized. It would also be of interest to further investigate the suggestion of Waller et al. (1997b) that optical Hubble type may be correlated with FUV--$V$ color (although rest--frame $V$--band data will not necessarily be available for high--redshift galaxies). If relations between UV morphology and physical properties can be found in local galaxies {\\it and} shown to extend to large look--back times, they could provide a valuable tool for studying the distant universe through rest--frame UV imaging."
    },
    "0002/astro-ph0002327_arXiv.txt": {
        "abstract": "{\\it Ginga} and Rossi X-ray Timing Explorer (RXTE) observations have allowed an unprecedented view of the recurrent systematic pulse shape changes associated with the 35-day cycle of Hercules X-1, a phenomenon currently unique among the known accretion-powered pulsars. We present observations of the pulse shape evolution. An explanation for the pulse evolution in terms of a freely precessing neutron star is reviewed and shown to have several major difficulties in explaining the observed pulse evolution pattern. Instead, we propose a phenomenological model for the pulse evolution based upon an occultation of the pulse emitting region by the tilted, inner edge of a precessing accretion disk. The systematic and repeating pulse shape changes require a {\\it resolved} occultation of the pulse emission region. The observed pulse profile motivates the need for a pulsar beam consisting of a composite coaxial pencil and fan beam but the observed evolution pattern requires the fan beam to be focused around the neutron star and beamed in the antipodal direction. The spectral hardness of the pencil beam component suggests an origin at the magnetic polar cap, with the relatively softer fan beam emission produced by backscattering from within the accretion column, qualitatively consistent with several theoretical models for X-ray emission from the accretion column of an accreting neutron star. ",
        "introduction": "Her X-1 is a 1.24 second period accretion-powered X-ray pulsar in a 1.7 day circular orbit with a normal stellar companion, HZ Her. In addition to these basic periodicities this binary system has long been known to display an unusual 35-day long cycle of High and Low X-ray flux states. Within a single 35-day cycle are found a Main High and Short High X-ray flux state lasting roughly ten and five days each respectively and separated by ten day long Low states (see e.g. Scott \\& Leahy 1999). Pulsations are detected during the High states but cease during the intervening Low states. The High states are punctuated by deep X-ray eclipses every 1.7 days indicating a line-of-sight close to the binary plane.  A phenomenon currently known to exist only in Her X-1 is a repeating, systematic evolution of the pulse profile that occurs during the 35-day cycle. Observations with the Large Area Counters on {\\it Ginga} and the Proportional Counter Array (PCA) on RXTE have allowed an unprecedented view of the evolution of the pulse profile in shape and energy spectrum across both the Main and the Short High states. The {\\it Ginga} observations cover the energy range 1-37 keV and sample five Main and two Short High states. They are described in detail in Deeter et al. (1991), Scott (1993) and Deeter et al. (1998). The bulk of the data consists of a Main-Short-Main High state sequence in April-May-June of 1989. Lightcurves and softness ratios for the 1988 and 1989 observations as well as some pulse profiles can be found in Leahy (1995a). The RXTE observations cover a consecutive Main and Short High state in September-October of 1996 and a Short-Main-Short High state sequence in September-October 1997. The 1996 Observations with the PCA in the range 2-60 keV are presented in Scott et al. (1997a).  In this paper, we present a simple phenomenological model for the pulse evolution based upon the occultation of the central X-ray source by the inner edge of a tilted, precessing, accretion disk. This choice is motivated by the observed association of pulse shape changes with {\\it decreases} in overall X-ray flux near the end of the Main High state (see Scott 1993; Scott et al. 1997a; Deeter et al. 1998; Joss et al. 1978 and Soong et al. 1990a) and the well known ability of a tilted, twisted, counter-precessing accretion disk to phenomenologically account for much of the complex 35-day optical and X-ray behavior displayed by the Her X-1/HZ Her system (Petterson 1975; Petterson 1977; Gerend \\& Boynton 1976; Crosa \\& Boynton 1980; Boynton, Crosa \\& Deeter 1980; Middleditch 1983). We stress that most of the observed 35-day behavior has been associated with the {\\it outer} portion of the accretion disk whereas we will demonstrate that the pulse evolution must be an {\\it inner} disk phenomenon.  Previous attempts to model the pulse shape evolution have relied on a combination of neutron star free precession and obscuration by the accretion disk (Tr\\\"umper et. al. 1986; Kahabka 1987, 1989), obscuration by ``flaps'' of matter at the juncture of the accretion disk and pulsar magnetosphere (Petterson et al. 1991) and obscuration of the pulsar by a tilted precessing disk (Bai 1981; Averitsev et al. 1992). None of these previous models attempted to explain more than a few aspects of the pulse evolution. The model presented here refines the disk and pulsar beam geometry to qualitatively account for the observed pulse shape and its evolution during both the Main and Short High states.  A discussion of relevant aspects of the tilted, twisted, counter-precessing disk is presented in section 2. A summary of the main features observed in the pulse evolution during the 35-day cycle is presented in section 3. Section 4 discusses the absolute pulse phase alignment of the Main and Short High state pulses. Section 5 discusses the pulse evolution as a consequence of neutron free precession. Section 6 briefly discusses pulse evolution as a consequence of changing mass accretion patterns onto the neutron star. Section 7 presents a pulse evolution model based on an inner disk occultation and a simple X-ray beam configuration. In section 8 we discuss some implications of this inner disk occultation interpretation. ",
        "conclusions": "The model proposed here for the pulse evolution cycle in Her X-1 consists of an occultation of the neutron star by the inner edge of a tilted and precessing disk. The leading edge decay of the main and interpulses during the Main High state preclude a beam geometry consisting of a pencil beam surrounded by a fan beam and emitting from the neutron star surface. However, reversing the fan beam so that is emitted in the antipodal direction with respect to the pencil beam and at some distance above the stellar surface allows the leading edge decay to be reproduced by an occultation in a natural fashion. The observed Short High state evolution pattern then arises naturally with this geometry.  Cyclotron resonant scattering in the accretion column is an attractive mechanism for producing a pencil beam surrounded by an reversed fan beam. The energy of the cyclotron resonance is directly proportional to the local magnetic field strength and this decreases with altitude above the neutron star surface. When the energy of an upward traveling photon equals the local cyclotron energy, scattering will occur. This creates a natural energy dependent filtering process for photons emitted in the accretion cap and divides the beam into three components. Hard photons (especially those above the surface cyclotron frequency) will escape in a pencil beam. Softer photons will be backscattered and gravitationally focused around the neutron star in an antipodally directed fan beam. Finally, photons scattered from the accretion column isotropically will produce a constant pulse component. The softest photons will scatter from the highest altitudes in the accretion column. The neutron star rotation will cause the highest altitude emission (D and E in figure 8) to rotate with pulse phase. The occultation of this high altitude emission by the precessing disk edge creates gaps in the constant emission profile and therefore a ``quasi-sinusoidal'' profile. A hard pencil beam (C), a softer fan beam (Ba and Bb) and a quasi-sinusoidal component (D and E) qualitatively match the observed energy dependence in the Main High state profile. This type of beam model may be applicable to other X-ray pulsars as well, for example Vela X-1, which displays a pulse profile consisting of two doubled peaked components at soft X-ray energies, possibly superposed on a quasi-sinusoidal component, that fills in at harder X-rays in a way quite similar to the Her X-1 main pulse (White, Swank \\& Holt 1983). In the case of Vela X-1, the observer's line-of-sight would have to be located much closer to the neutron star spin equator than in Her X-1.  While an occultation model apparently has many attractive features, a roughly $45\\deg$ tilt is required between the neutron star spin axis and the binary axis of the system. Is this plausible? A large tilt to the neutron star may have originated in the supernova explosion that gave it birth. The accretion of matter from the companion will cause angular momentum to be accreted by the neutron star. Since the direction of the time averaged accreted angular momentum is along the binary axis of the system, the neutron star's spin axis should become coaligned with the binary axis. This event will take some time and the accreted angular momentum will spin up the neutron star. In fact, the alignment timescale and the spin-up timescale should be comparable. The measured spin-up time of Her X-1 is about $10^{5}$ years. This timescale is the same as that predicted for the entire X-ray emitting phase of Her X-1 (Savonije 1978). Therefore, if Her X-1 is currently tilted that tilt will in all likelyhood be maintained throughout the rest of the X-ray emitting phase. However the important question is what is the {\\it current} ratio of accreted angular momentum to that at the start of the X-ray emitting phase?  Her X-1 may be near the start of the X-ray emitting phase. The X-ray emitting phase ends when the accretion flow becomes great enough to smother the neutron star and prevent pulsations. Historical optical observations of HZ Her show that the X-ray heating and hence the mass accretion rate has ceased occasionally for years to decade-long periods over the last hundred years (Jones, Forman \\& Liller 1973; Hudec \\& Wenzel 1986). The current state of mass transfer depends on the X-ray heating of HZ Her and is inhibited by the X-radiation pressure. Thus HZ Her must be close to, but not quite, filling its Roche lobe. This state of affairs is more consistent with Her X-1 being at the beginning rather than the end of its X-ray phase. If so, then a highly tilted neutron star can plausibly exist in the Her X-1 system. Another point in favor of such an interpretation is the 1.24 second pulse period of Her X-1. This period is typical of the radio pulsar population. A period longer than 3 seconds or much shorter than 1 second would unequivocally show that Her X-1 has been spun down or spun up by a significant history of interaction with circumstellar matter. Finally, a tilted neutron star should cause a persistent asymmetry in the optical orbital lightcurve about orbital phase 0.5. No such asymmetry was reported by Deeter et al. (1976) but later studies of the orbital optical lightcurve with more data indicate the presence of just such an asymmetry (e.g. Voloshina, Lyuti \\& Sheffer 1990; Thomas et al. 1983).  The properties required of the inner disk to fit the observed pulse evolution are a scale height comparable to the neutron star diameter and a small inner radius ($20-40 R_{ns}$). For comparison, the corotation radius is $157R_{ns}$ for Her X-1. The predicted magnetospheric radius, $R_m$, for disk accretion depends on assumptions on the boundary conditions. Using the known parameters for Her X-1 (Leahy \\& Scott 1998), one obtains $R_m=3.5\\times10^8\\alpha^{-2/61}$ cm for the model of Kiraly \\& M\\'esz\\'aros (1988); $R_m=3.8\\times10^8\\gamma^{2/7}$ cm for the model of Lamb (1988); and $R_m=4.9\\times10^8$ K cm for the formula of Finger et al. (1996) (which has the special case of $K=0.47$ for the model of Ghosh \\& Lamb, 1978). The model of Aly (1980) for a highly conducting disk, gives only a slightly smaller value for the disk inner radius: $R_m=3.0\\times10^8\\alpha^{2/7}\\cos(\\chi)^{4/7}$ cm (for $\\chi$ near $\\pi/2$, the dependence on $\\chi$ is $\\sin(\\chi)^{40/69}$, so the minimum $R_m$ is $\\simeq1/3$ of this for $\\chi=0$). The disk viscosity parameter is $\\alpha$, $\\gamma$ is defined in Lamb (1988), K is a dimensionless parameter, which Finger et al. (1996) find to be $\\simeq 1$ for A0535+26, and $\\chi$ is the tilt of the dipole axis from the equatorial plane. In all cases the predicted disk radius greatly exceeds the inner disk radius required by the occultation model.  Another observation relevant to the magnetospheric radius is the spinup rate of Her X-~1. It is the smallest among the accretion-powered pulsars and indicates the slowest rate of net angular momentum accretion. During the giant outbursts of Be-transients, an X-ray flux vs. spin-up rate correlation has been observed in which spin-up occurs at a rate consistent with the fiducial torque (e.g. Finger et al. (1996); Bildsten et al. (1997); Nelson et al. (1997); Scott et al. (1997)). The formula for the fiducial accretion torque $N_f$ is given by $N_f=\\dot M(GM_{X}R_s)^{1/2}$ where $R_s$ is the magnetospheric radius for spin-up, $\\dot M$ is the mass accretion rate and $M_{X}$ is the mass of the neutron star. This formula assumes all the angular momentum of the accreting matter at $R_s$ is given to the neutron star. The known spin-up rate gives a value for $R_s=8.5$ km: clearly too small to be physical. The small disk radius from the occultation model is much larger than this. This implies that the actual torque on the neutron star must be smaller than the fiducial torque, to allow a magnetospheric radius as large as $20-40R_{ns}$. In the Ghosh \\& Lamb (1978, 1979) model of an aligned rotator, a residual portion of the magnetic field is not screened by currents near the magnetospheric radius and will interact with material orbiting beyond the corotation radius. The interaction adds a negative torque component to the total torque. A smaller net torque will be exerted if the magnetospheric radius approaches the corotation radius\\footnote{Note that Ghosh and Lamb prefer to call the corotation radius the `centrifugal radius' and the magnetospheric radius the `corotation radius'.} of the neutron star. Thus nearly any spin-up rate can be obtained within a small range of magnetospheric radii near to but inside the corotation radius. So it is not surprising that the predicted magnetospheric radius from the Ghosh and Lamb model for Her X-1 is nearly equal to the corotation radius given above. However, it is well known that the Ghosh and Lamb model is not consistent with the spin behaviour of many X-ray pulsars (e.g. Bildsten et al. 1997; Nelson et al. 1997). In summary, different models give a magnetospheric radius from the spin-up rate of Her X-1 in a wide range between $R_{ns}$ and $R_c=157R_{ns}$.  How does one account for an inner disk edge which is at $20-40R_{ns}$? It is much smaller than the predicted neutron star magnetosphere radius, yet may be much larger than the radius predicted from simple spin-up. We speculate on what may be the real physical origin of the $20-40R_{ns}$ disk inner edge. One possibility is that the dipole magnetic field is much smaller than that indicated by the observed X-ray cyclotron line so that the predicted magnetospheric radius is much smaller. The magnetospheric radius decreases from $\\sim 200R_{ns}$ to $30R_{ns}$ if the dipole component of B is a factor of 25 less than deduced from the cyclotron line. The cyclotron line would be explained as arising in non-dipolar magnetic fields in the accretion region at the neutron star surface. An emission region concentrated in a non-dipolar surface pocket of the field and producing a pencil beam would likely be difficult to distinguish from a similar emission region in the case of a pure dipole. Another explanation has been put forth by Baushev \\& Bisnovatyi-Kogan (1999) in which a magnetic field of $4-6 \\times 10^{10}$ G is estimated for Her X-1 from the observed cyclotron line energy by assuming a large anisotropy exists in the electron momentum distributions parallel and perpendicular to the magnetic field lines. Whether either of these possibilities is viable will probably require much more research.  An alternative to a reduced dipole field is a stable, thin inner disk which penetrates far into the magnetosphere. However, this seems very unlikely: the magnetosphere is very stiff due to the steep gradient of magnetic pressure and energy density (as $r^{-6}$) for a dipole field. Models which have the disk penetrate as deeply as possible (e.g. Aly 1980), do not have an inner radius very much less than other models (tilting the magnetic axis gives a reduction at most by a factor 3 in the inner disk radius). So a disk penetrating far into the magnetosphere appears to be unfeasible.  Another possibility is that current models for determining the magnetospheric radius are inadequate. For example, Miller \\& Stone (1997) use magnetohydrodynamic calculations to show that the Balbus-Hawley instability  and magnetic braking have dramatic effects on the magnetospheric boundary. They also result in outflowing winds along field lines opened up by reconnection. It is quite possible that the small inner disk in Her X-1 is due to such effects. We note that mass outflows are quite likely as there is extended X-ray emission from a large corona in Her X-1.  The model described in this paper is phenomenological and highly idealized. The purpose was to show that an inner disk occultation can explain the observations and has reasonable physical grounds for support. Considerable refinement of the model remains to be done. The decay of the soft shoulders of the main pulse occurs earlier in the Main High state than an axisymmetric beam model predicts. The soft shoulders of the Main High state main pulse are unequal in amplitude. These problems might indicate an offset in the dipole axis.  An occultation of the neutron star by the inner accretion disk explains the pulse evolution cycle of Her X-1 in a natural fashion. Many of the details of the observed evolution can be accounted for by invoking a reversed fan beam geometry around a neutron star significantly inclined to the binary axis of the system. The leading edge decay of the main pulse and interpulse during the Main High state is properly predicted. During the Short High state, the rapid disappearance of the small hard peak as well as the slower decay of the soft peak are predicted. Note that the soft peak shows little decay of either the leading or trailing edge and this is also predicted by the occultation model. In summary, most of the curious pulse shape changes observed during the Main and Short High states are tied together with a single simple occultation model. We note that no previous model for the pulse shape changes, including free precession, has been able to account for the observed details of the pulse evolution."
    },
    "0002/astro-ph0002057_arXiv.txt": {
        "abstract": "After studying the orientation of a warp generated by a companion satellite, we show that the Galactic Warp would be oriented in a different way if the Magellanic Clouds were its cause. We have treated the problem analytically, and complemented it with numerical N-Body simulations. ",
        "introduction": "The disk of the Milky Way is remarkably flat out to 10 kpc, where it starts to bends in opposite directions in the southern and northern parts. The cause of it is still a puzzle: for a review, see Binney \\shortcite{b92}.  One possibility is that the Magellanic Clouds distort the Galactic disk in the observed way. The fact that the direction of the maximum warping lies very close to the galactocentric longitude of the Clouds makes this hypothesis tempting.  The problem with this scenario is that the Clouds are not massive enough to generate the warp amplitudes that we observe at their present distance. This was noticed by the first researches in this field \\cite{bu57,ke57}, and later by Hunter \\& Toomre (1969). A remedy which might allow this scenario to work was to suppose that the Clouds are currently not at the pericenter of their orbit, so that they have been much closer to the Galactic disk in a previous epoch ($\\simeq 20$ kpc is what was needed). However, later work by Murai \\& Fujimoto (1986) determined the orbit of the Clouds, and proved that the Clouds are actually at their pericenter, with an apocenter close to 100 kpc, so the problem of the small amplitude still remains if the Clouds are to be blamed as the cause of the Galactic warp.  Recently a mechanism for amplifying the effect of a satellite has been proposed by Weinberg \\shortcite{w98}. He describes a calculation in which a disk galaxy surrounded by a dark halo is perturbed by a massive satellite, similar to the LMC. By means of a linear perturbation analysis, he follows the perturbation (wake) created by the satellite in the halo, including its self-gravity. He finds that the torque exerted on the disk is several times larger than that due directly to the satellite: the latter is amplified because (i) the satellite-induced wake in the halo itself exerts a torque, roughly in phase with that from the satellite; and (ii) the wake itself further perturbs the halo, resulting in a torque that is larger again. Under circumstances in which the satellite orbital frequency is close to the natural precession frequency of the disk (i.e., the warp mode frequency of Sparke \\& Casertano \\shortcite{sc}), a significant amplitude can result.  A calculation by Lynden-Bell \\shortcite{lb85} of a similar scenario gives comparable results, as does a simple model described by Kuijken \\shortcite{k97}.  In this paper, we focus on the orientation of a warp generated by a massive orbiting satellite with less emphasis on the amplitude of the warp. In \\S2 we discuss a simple analytic model in which the disk and halo are rigid: this establishes the baseline response of a disk to satellite tides, and its orientation with respect to the satellite orbital plane. As we show, this orientation is different from that of the Galactic warp to the LMC orbit. \\S3 contains a description of the N-body code used, and \\S4 to \\S6 the results of our N-body simulations, showing that the orientation problem remains. In \\S7 we give our conclusions. ",
        "conclusions": "We show by means of analytic work and N-body simulations that the precession path of a warp generate by an orbiting satellite galaxy is elongated along the direction perpendicular to the satellite's orbital plane. Applying our result to the Milky Way, if the Galactic Warp is generated by the Magellanic Clouds, the direction of maximum amplitude of the warp would lie close to $l \\simeq 0^\\circ$, as compared to the observed direction of $\\simeq 90^\\circ$. Even if the halo's self-gravitating tidal response to the satellite amplifies the effect of the satellite~\\cite{w98}, this response will be mostly in phase with the satellite, and the alignment problems will persist. Possibly the Sagittarius dwarf galaxy, whose orbit lies at 90$^\\circ$ to that of the LMC, can be the cause of the warp instead~\\cite{ibata}.  A limitation of the present work is that the halo has not been considered as a live, self-gravitating component. It has been shown~\\cite{dk,nt} that the back-reaction of the halo on a re-aligning disk can have important consequences. Such effects will be the subject of a further paper.  \\appendix"
    },
    "0002/astro-ph0002261_arXiv.txt": {
        "abstract": "We present a detailed calculation of the evolution of low--mass ($< 0.25~M_\\odot $) helium white dwarfs. These white dwarfs (the optical companions to binary millisecond pulsars) are formed via long--term, low--mass binary evolution. After detachment from the Roche lobe, the hot helium cores have a rather thick hydrogen layer with mass between 0.01 to 0.06$~M_\\odot $. Due to mixing between the core and outer envelope, the surface hydrogen content is 0.5 to 0.35, depending on the initial value of the heavy element (Z) and the initial secondary mass. We found that the majority of our computed models experience one or two hydrogen shell flashes. We found that the mass of the helium dwarf in which the hydrogen shell flash occurs depends on the chemical composition. The minimum helium white dwarf mass in which a hydrogen flash takes place is 0.213$~M_\\odot $ (Z=0.003), 0.198$~M_\\odot $ (Z=0.01), 0.192$~M_\\odot $ (Z=0.02) or 0.183$~M_\\odot $ (Z=0.03). The duration of the flashes (independent of chemical composition) is between few $\\times 10^6 $ years to few $\\times 10^7 $ years. In several flashes the white dwarf radius will increase so much that it forces the model to fill its Roche lobe again. Our calculations show that cooling history of the helium white dwarf depends dramatically on the thickness of the hydrogen layer. We show that the transition from a cooling white dwarf with a temporary stable hydrogen--burning shell to a cooling white dwarf in which almost all residual hydrogen is lost in a few thermal flashes (via Roche--lobe overflow) occurs between 0.183--0.213$~M_\\odot $ (depending on the heavy element value). ",
        "introduction": "Kippenhahn, Kohl \\& Weigert (1967) were the first who followed the formation of helium white dwarfs (WD) of low mass in a binary system. The evolution of a helium WD of 0.26$~M_\\odot$ (remnant) was investigated by Kippenhahn, Thomas \\& Weigert (1968) who found that a hydrogen flash can be initiated near the base of the hydrogen rich envelope. The energy of the flash is sufficient to cause the envelope to expand to giant dimensions and hence it may be possible that another short term Roche lobe filling can occur.  In Webbink (1975), models of a helium white dwarf were constructed by formally evolving a model from the homogeneous zero--age main sequence with the reduction of the mass of the hydrogen--rich envelope. When the mass of the envelope is less than some critical value, the model contracts adopting white dwarf dimensions. Webbink found that thermal flashes do not occur for WDs less massive than 0.2$~M_\\odot$. Alberts et al. (1996) have confirmed Webbink's finding that low--mass white dwarfs do not show thermal flashes and the cooling age for WDs of mass $M_{wd}$$\\leq$0.20$M_\\odot$ can be considerably underestimated if using the traditional WD cooling curves which were constructed for $ M_{wd}$$>$0.3$M_\\odot$ (Iben \\& Tutukov 1986, IT 86).  Recently, Hansen \\& Phinney (1998a -- HP98) and Benvenuto \\& Althaus (1998 -- BA98) investigated the effect of different mass of the hydrogen layer ($\\rm 10^{-8} \\le M_{env}/M_\\odot \\le 4 \\times 10^{-3} $) on the cooling evolution of $\\rm 0.15 \\le M_{He}/M_\\odot \\le 0.5 $ helium WDs. In both calculations (BA98 and HP98) the mass of the hydrogen envelope left on the top of white dwarf has been taken as free parameter. BA98 found that thick envelopes appreciably modify the radii and surface gravities of no--H models, especially in the case of low--mass helium white dwarfs.  Driebe et al. (1998 -- DSBH98) present a grid of evolutionary tracks for low-mass white dwarfs with helium cores in the mass range from 0.179 to 0.414$~M_\\odot$. The tracks are based on a 1$~M_\\odot$ model sequence extending from the pre--main sequence stage up to the tip of red giant branch. Applying large mass loss rates forced the models to move off the giant branch and evolve across the Hertzsprung--Russell diagram and down the cooling branch. They found that hydrogen flashes take place only for two model sequences, 0.234$~M_\\odot$ and 0.259$~M\\odot$, and for very low--mass WDs the hydrogen shell burning remains dominant even down to effective temperatures well below 10 000 K. According to our previous calculations (Ergma, Sarna \\& Antipova, 1998) we  find that for a low--mass white dwarf with a helium core, which was formed during low--mass binary evolution (after detachment from the Roche lobe), the hydrogen layer left on the top of the helium core is much thicker ($\\rm \\sim 1-6 \\times 10^{-2} M_\\odot $ with $\\rm X_{surf} $ ranging from 0.3 to 0.52) than used in cooling calculation by HP98 and BA98. Also in DSBH98 (see their Table 1), for the two lowest total remnant masses the envelope mass value is smaller that obtained in our calculations. ",
        "conclusions": "We have performed comprehensive evolutionary calculations to produce a close binary system consisting of a NS and a low--mass helium WD.  We argue that the presence of a thick hydrogen layer changes dramatically the cooling time--scale of the helium white dwarf ($\\rm < 0.25 ~M_\\odot $), compared to the  previous calculations (HP98, BA98) where the mass of the hydrogen envelope was chosen as free parameter and was usually one order of magnitude less than that obtained from real binary evolution computations.  Also, we have demonstrated that using new cooling tracks we can consistently explain the evolutionary status of the binary pulsar PSR J1012+53.  Tables with cooling curves are available on {\\bf http://www.camk.edu.pl/~sarna/}."
    },
    "0002/astro-ph0002507_arXiv.txt": {
        "abstract": "We present a study of 81 $I$-band selected, spectroscopically-confirmed members of the X-ray cluster \\1054{} at $z=0.83$. Redshifts and spectral types were determined from Keck spectroscopy.  Morphologies and accurate colors were determined from a large mosaic of HST WFPC2 images in $R_{F606W}$ and $I_{F814W}$, corresponding to $U$ and $B$ in the restframe.  Early-type galaxies constitute only $44$\\,\\% of this galaxy population. This fraction is much lower than in comparable rich clusters at low redshift. Thirty-nine percent are spiral galaxies, and 17\\,\\% are mergers.  The early-type galaxies follow a tight and well-defined color-magnitude relation, with the exception of a few outliers.  The observed scatter is $0.029 \\pm 0.005$ magnitudes in restframe $U-B$.  Most of the mergers lie close to the CM relation defined by the early-type galaxies. They are bluer by only $0.07 \\pm 0.02$ magnitudes, and the scatter in their colors is $0.07 \\pm 0.04$ magnitudes. Spiral galaxies in \\1054{} exhibit a large range in their colors.  The bluest spiral galaxies are $\\sim 0.7$ magnitudes bluer than the early-type galaxies, but the majority is within $\\pm 0.2$ magnitudes of the early-type galaxy sequence. The red colors of the mergers and the majority of the spiral galaxies are reflected in the fairly low Butcher-Oemler blue fraction of \\1054{}: $f_B=0.22 \\pm 0.05$, similar to intermediate redshift clusters and much lower than previously reported values for clusters at $z \\sim 0.8$. The slope and scatter of the CM relation of early-type galaxies are roughly constant with redshift, confirming previous studies that were based on ground-based color measurements and very limited membership information.  However, the scatter in the combined sample of early-type galaxies and mergers (i.e., the sample of {\\em future} early-type galaxies) is twice as high as the scatter of the early-type galaxies alone. This is a direct demonstration of the ``progenitor bias'': high redshift early-type galaxies seem to form a homogeneous, old population because the progenitors of the youngest present-day early-type galaxies are not included in the sample. ",
        "introduction": "The star formation history of early-type galaxies provides an important constraint on models for galaxy formation and a test for stellar population synthesis models. Because present-day early-type galaxies consist mostly of old stars, the evolution of their stellar populations may provide clues to conditions in the early Universe. Early-type galaxies (elliptical galaxies and S0 galaxies) are very efficiently studied in clusters.  They form the dominant population in nearby rich clusters (Oemler 1974; Dressler 1980), and most of our knowledge of the formation and evolution of early-type galaxies has come from studies of clusters at $0<z<0.5$.  Strong constraints on the star formation histories of early-type galaxies have been derived from the redshift evolution of the Fundamental Plane relation (van Dokkum \\& Franx 1996, Kelson et al.\\ 1997, Pahre 1998, Bender et al.\\ 1998, van Dokkum et al.\\ 1998b, J\\o{}rgensen et al.\\ 1999, Kelson et al.\\ 2000) and the color-magnitude relation  (e.g., Ellis et al.\\ 1997, Stanford, Eisenhardt, \\& Dickinson 1998, van Dokkum et al.\\ 1998a, Kodama et al.\\ 1998). The strongest constraints on the {\\em mean} star formation epoch are provided by the evolution of the Fundamental Plane, mainly because the mean luminosity evolution is large compared to the mean color evolution. van Dokkum et al.\\ (1998b) find a slow luminosity evolution of early-type galaxies to $z=0.83$, and conclude that their stars were formed at $z>2.8$ for a Salpeter (1955) IMF, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0$, or $z>1.7$ for $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$.  Whereas the Fundamental Plane can be used to measure the mean luminosity evolution of early-type galaxies with high accuracy, it is difficult to obtain a sufficiently large sample to study the evolution of its slope and scatter (but not impossible, see Kelson et al.\\ 2000). Furthermore, the Fundamental Plane can only be applied to early-type galaxies, and concentrating on the early-type galaxy population alone may lead to an incomplete picture of their evolution. There is growing evidence that a significant fraction of present-day early-type galaxies was assembled relatively recently (Dressler et al.\\ 1997; van Dokkum et al.\\ 1999).  They may have formed from gas stripping of spiral galaxies (e.g., Gunn \\& Gott 1972, Tytler \\& Vidal 1978, Solanes \\& Salvador-Sole 1992, Couch et al.\\ 1998, Abadi, Moore, \\& Bower 1999), through tidal interactions (e.g., Fried 1988, Lavery, Pierce, \\& McClure 1992, Oemler, Dressler, \\& Butcher 1997, Moore, Lake, \\& Katz 1998, Couch et al.\\ 1998), or in mergers (e.g., Lavery et al.\\ 1992, van Dokkum et al.\\ 1999). Early-type galaxies in high redshift clusters may therefore form a special subset of the progenitors of early-type galaxies in nearby clusters.  Failure to consider the other progenitors may result in biased age estimates (see, e.g., Franx \\& van Dokkum 1996).  The evolution of the color-magnitude (CM) relation can provide important constraints additional to those provided by the Fundamental Plane. The CM relation is useful for studying all morphological types and large photometric samples can be obtained straightforwardly. The combination of the mean color evolution of galaxies and their luminosity evolution can constrain the IMF (e.g., Kelson et al.\\ 2000). The evolution of the slope of the CM relation is sensitive to age differences between high mass galaxies and low mass galaxies (see, e.g., Kodama et al.\\ 1998). The scatter in the CM relation is particularly important. It can be measured to high accuracy with HST and it constrains the age spread of galaxies of a given luminosity (Bower, Lucey, \\& Ellis 1992; Ellis et al.\\ 1997; van Dokkum et al.\\ 1998a; Stanford et al.\\ 1998). The scatter is proportional to $\\Delta \\ln (\\tau) = \\Delta \\tau / \\langle \\tau \\rangle$, where $\\tau$ is the luminosity-weighted age (see van Dokkum et al.\\ 1998a).  The evolution of the color-magnitude relation has been studied extensively (e.g., Aragon-Salamanca et al.\\ 1993, Stanford, Eisenhardt, \\& Dickinson 1995, Ellis et al.\\ 1997, Stanford et al.\\ 1998, Kodama et al.\\ 1998). These studies are in remarkable agreement: the slope and scatter seem to be roughly constant with redshift, and the mean color evolution is consistent with passive evolution of an old stellar population that was formed at high redshift. However, these studies suffer from limited membership information and areal coverage (usually one HST+WFPC2 pointing). Because the field contamination among late-type galaxies becomes very large at $z \\sim 1$ usually only early-type galaxies are considered, and even these morphologically selected samples suffer from significant contamination. For example, Stanford et al.\\ (1998) estimate that the field contamination among the early-type galaxies in their $z>0.6$ clusters is $20 - 40$\\,\\%, depending on the richness of the cluster. An additional complication is that the true contamination is expected to vary considerably on the small angular scales of the fields used for these studies.  We are undertaking a program to obtain large field HST+WFPC2 mosaics of clusters at $0.3 < z < 1$, combined with extensive spectroscopy from the ground. This program is complimentary to other HST studies of distant clusters, in the sense that we observe only a few clusters, but over a large field and with redshifts for $\\gtrsim 200$ (field and cluster) galaxies for each cluster field. The CM relation in the cluster CL\\,1358+62 at $z=0.33$ was presented by van Dokkum et al.\\ (1998a). We used a sample of 194 confirmed members observed with HST. It was found that S0 galaxies in the outskirts of the cluster show a much larger scatter in their colors and are bluer on average than those in the central regions, providing evidence for recent infall of galaxies from the field. Here, we present extensive Keck spectroscopy and a large HST WFPC2 mosaic in two filters of the cluster \\1054{} at $z=0.83$. Redshifts, morphologies and accurate colors are presented for a sample of 81 confirmed cluster members. One of the most striking results of our survey is the large number of merging systems in \\1054{} (van Dokkum et al.\\ 1999). In the present study we focus on the CM relation of confirmed members. The aims are to test whether the early-types are a homogeneous population at $z \\approx 1$, and to compare their CM relation to those of other morphological types. ",
        "conclusions": "We have presented a study of the galaxy population in the rich cluster \\1054{} at $z=0.83$. Our sample of 81 $I$-band selected galaxies is currently the largest sample of confirmed cluster galaxies at $z>0.5$ observed with HST.  One of the most striking results of our survey of \\1054{} is the high fraction of mergers among the $L>L_*$ cluster population (van Dokkum et al.\\ 1999). In the present paper we focussed on the early-type fraction, the Butcher-Oemler effect, and the color-magnitude relation in \\1054.  The early-type fraction in \\1054{} is much lower than that in rich clusters at intermediate and low redshift. Combining our data with data from the literature we find a continuous decrease in the early-type fraction with redshift. The early-type fraction changes by a factor $\\sim 2$ from $z=0$ to $z=0.83$. This result is robust and independent of the classification procedure. It is reasonable to assume that the mergers in \\1054{} will evolve into early-type galaxies. This implies that the early-type fraction in this cluster will increase from $44$\\,\\% to 61\\,\\% in approximately 1\\,Gyr.  \\1054{} displays a similar Butcher-Oemler effect as clusters at intermediate redshift: $f_B = 0.22 \\pm 0.05$. This may appear to be at variance with the low early-type fraction, but is a result of the fact that most spiral galaxies are red, similar to spiral galaxies in nearby clusters (e.g., Butcher \\& Oemler 1984). Of eleven spiral galaxies within $R=R_{30}$, only three are blue. The blue fraction in \\1054{} is much lower than values reported by Rakos \\& Schombert (1995) for clusters at $z>0.6$. This is probably due to the fact that these authors used a different procedure to measure the blue fraction.  We studied the CM relation as a function of morphology. The early-type galaxies follow a remarkably tight and well defined relation, with an observed scatter of $0.029 \\pm 0.005$ in restframe $U-B$. There is no significant dependence of the CM relation of early-type galaxies on radius in the cluster; our small sample precludes us from confirming or ruling out a similar trend as found for S0 galaxies in the cluster CL\\,1358+62 at $z=0.33$ (van Dokkum et al.\\ 1998a). The spiral galaxies exhibit a large range in their colors; they are offset to the blue by $0.22$ magnitudes with respect to the CM relation of the early-type galaxies, but many are as red as early-type galaxies of the same luminosity. This could be due to a spread in age, metallicity, and/or dust content.  The mergers are only slightly bluer than the early-type galaxies, by $0.07 \\pm 0.02$ magnitudes. The scatter in their CM relation is $0.07 \\pm 0.04$. Only three of the thirteen mergers are more than 0.2 magnitudes bluer than the CM relation; these are the only mergers with detected [O\\,{\\sc ii}] 3727\\,\\AA{} emission. It may appear surprising that the mergers are so red; mergers in the nearby Universe are typically blue and have high star formation rates (Liu \\& Kennicutt 1995). However, most of the spiral galaxies in \\1054{} are red as well, and the red colors of the mergers are probably a result of the environment: galaxies in and near a rich cluster are likely redder than typical field galaxies. A quantitative demonstration of this effect is that the median color of the mergers ($\\langle U-B \\rangle_z = 0.36$) is very similar to the median color of the full sample of cluster galaxies ($\\langle U-B \\rangle_z = 0.41$).  The slope and scatter of the CM relation of the early-type galaxies in \\1054{} are very similar to those of clusters at lower redshift, confirming previous studies. However, the presence of the mergers and the low early-type fraction in \\1054{} demonstrate that early-type galaxies in distant clusters are a biased set, and not representative of the full sample of progenitors of present-day early-type galaxies. This ``progenitor bias'' has a large effect on the scatter in the color-magnitude relation: the scatter in the full sample of all progenitors of present-day early-type galaxies (i.e., including the mergers in \\1054{}) shows strong evolution with redshift. The scatter increases by a factor $\\sim 2$ between $z=0$ and $z=1$. The main uncertainty in this result is the assumption that the galaxy population in \\1054{} is typical for its redshift.  It is clearly important to extend this type of analysis to other clusters, and to higher redshift.  Of particular interest  is the evolution of the early-type fraction, since it may be a fairly robust measure of morphological evolution. The large format of the HST imaging was very important in our study, because the mergers typically reside at large radii (van Dokkum et al.\\ 1999), and it allowed efficient spectroscopic follow-up. This study demonstrates that the combination of extensive spectroscopy and large format, high resolution imaging can provide important insights into the assembly and evolution of cluster galaxies."
    },
    "0002/astro-ph0002488_arXiv.txt": {
        "abstract": "We study the effect of microturbulence in the line formation calculations of H and He lines, in the parameter range typical for O and early B stars. We are specially interested in its effect on the determination of stellar parameters: $T_\\mathrm{eff}$, $\\log g$ and specially on the He abundance.  We first analyze the behaviour of H and He model lines between 4\\,000 and 5\\,000 \\AA~ with microturbulence and find that for O stars only \\ion{He}{i} lines and \\ion{He}{ii} $\\lambda$ 4686 are sensibly affected by microturbulence, and that models with lower gravities, the ones suitable for supergiants, are more sensitive to it.  Using a test procedure we show that the expected changes in $T_\\mathrm{eff}$, $\\log g$ and Helium abundance due to the inclusion of microturbulence in the analysis, are small. We analyze five stars (two late, one intermediate and two early O stars) using microturbulence velocities of 0 and 15 kms$^{-1}$ and confirm the result of the previous test. The parameters obtained for 15 kms$^{-1}$ differ from the ones at 0 kms$^{-1}$ within the limits of the standard error box of our analysis. Only later types reduce their He abundance, by 0.02 in $\\epsilon$. Comparing with values in the literature we find that the range of our changes agree with previous results. In some cases other effects can add to microturbulence, and further reduce the He abundance up to 0.04. The quality of the line fits only improves for \\ion{He}{i} $\\lambda$ 4471, but not to the extent of completely solving the so--called dilution effect.  Therefore our conclusion is that microturbulence is affecting the derivation of stellar parameters, but its effect is comparable to the adopted uncertainties. Thus it can reduce moderate He overabundances and solve line fit quality differences, but it cannot explain by itself large He overabundances in O stars. ",
        "introduction": "Although microturbulence is a fundamental parameter when deriving abundances from stellar spectra, it has been systematically ignored when deriving He abundances from quantitative spectroscopical NLTE analyses of OB stars (Herrero et al. \\cite{Herrero92}, Smith \\& Howarth \\cite{Smith&Howarth94}), being the first exceptions to this the works by McErlean et al. (\\cite{McErlean98}) and Smith \\& Howarth (\\cite{Smith&Howarth98}).  One of the main reasons to ignore it is that the use of NLTE techniques reduced the need of microturbulence for the reproduction of the observed metallic line strengths and they even made abundances less sensitive to the value adopted for microturbulence than in LTE analyses (Becker \\& Butler \\cite{Becker&Butler89}).  In quantitative NLTE analyses of OB stars, the lines of H and He are used to determine stellar parameters (namely, effective temperature, $T_\\mathrm{eff}$, logarithmic surface gravity, $\\log g$, and He abundance). Their line profiles are dominated by the Stark broadening, and microturbulence, included in the standard way, only adds an extra Doppler width to the thermal broadening. First values for microturbulence found in LTE for main sequence B stars, of the order of 5 kms$^{-1}$ (Hardorp \\& Scholz \\cite{Hardorp&Scholz70}), showed to be small enough to be of no importance when compared to H and He thermal velocities, that in such hot atmospheres are well above 5 kms$^{-1}$. This together with the use of NLTE made the influence of microturbulence in the overall profile seem negligible.  When measuring microturbulence in OB supergiants the situation is quite different, as its value has always showed to be comparable or well above the thermal velocities of H and He, and even, in some cases, above the speed of sound in these atmospheres (Lamers \\cite{Lamers72}, Lennon \\& Dufton \\cite{Lennon&Dufton86}). This was not changed by later NLTE analyses (Lennon et al. \\cite{Lennon91}, Hubeny et al. \\cite{Hubeny91}, Gies \\& Lambert \\cite{Gies&Lambert92}, Smartt et al. \\cite{Smartt97}), in which the values of microturbulence obtained are usually reduced but never to values clearly without contradiction. Kudritzki (\\cite{Kudritzki92}) and Lamers \\& Achmad (\\cite{Lamers&Achmad94}) explained that this could be due to the presence of wind outflow in these stars, that can mimic the effect of microturbulence in the line profiles.  Most of the previously referred works are on early B and late O supergiants. Little work has been done on early O stars (see Hubeny et al. \\cite{Hubeny91}), mainly because for them, with scarce metallic lines, it is very difficult to measure the value of microturbulence. And it is in the whole range of O stars that we are interested.  One problem that is systematically found in all analyses of OB spectra is the imposibility of finding a consistent fit of all \\ion{He}{i} lines with a unique set of parameters. This difficulty appears worse in supergiants than in main sequence stars, and it is bigger between results from singlet and from triplet lines, but also appears within different lines in one system. We call this the {\\emph {\\ion{He}{i} lines problem}}.  This discrepancy between singlet and triplet lines of \\ion{He}{i} was investigated by Voels et al. (\\cite{Voels89}), who considered that atmospheric extension was responsible and so they called this effect ``generalized dilution effect''. But recent works with spherical, mass losing models show that this can not be the only reason (Herrero et al., \\cite{Herrero95}, \\cite{Herrero00}). Furthermore, McErlean et al. (\\cite{McErlean98}) and Smith \\& Howarth (\\cite{Smith&Howarth98}) find that the inclusion of microturbulence in the line formation calculations reduces, although not to all its extent, the discrepancy between the fits of singlet and triplet lines. They also find better and more consistent fits for the whole set of \\ion{He}{i} lines when microturbulence is considered.  One more reason that encourages us to study the effect of microturbulence is that a careful inspection of several works shows that even for main sequence stars,values of microturbulence comparable to the thermal velocities of at least He are obtained in NLTE (Gies \\& Lambert \\cite{Gies&Lambert92}, Kilian et al. \\cite{Kilian91}), which can invalidate the hypothesis of negligible influence of microturbulence in the profiles.  So in this paper we want to study microturbulence in the range of parameters typical for O stars of any luminosity class, specifically pointing to its effect on the determination of stellar parameters and also on the {\\emph {\\ion{He}{i} lines problem}}. We are specially interested in investigating whether the inclusion of microturbulence can solve the {\\it He discrepancy} (Herrero et al. \\cite{Herrero92}), as Vrancken (\\cite{Vrancken97}), McErlean et al. (\\cite{McErlean98}) and Smith \\& Howarth (\\cite{Smith&Howarth98}) suggest. This now well known problem has induced a lot of theoretical improvements in both evolutionary and model atmosphere theories in order to explain it. Recent works including sphericity and mass loss in the spectral analysis (Herrero et al. \\cite{Herrero95}, \\cite{Herrero00}, Israelian et al. \\cite{Israelian99}) show that with these new models, more suitable for these stars than the plane--parallel hydrostatic ones, the discrepancy is not solved. Evolutionary models including additional mixing inside the star can explain enhanced He photospheric abundances during the H-burning phase, this mixing being induced by different physical mechanisms like rotation and turbulent diffusion (Dennisenkov, \\cite{denn94}, Meynet \\& Maeder, \\cite{Meynet&Maeder97}, Maeder \\& Zahn, \\cite{Maeder&Zahn98}), or turbulent diffusion and semiconvection (Langer \\cite{Langer92}). Sometimes angular rotational velocity changes during evolution are included (Langer \\& Heger, \\cite{LH98}, Heger, \\cite{Heger98}). For a recent review, see Maeder \\& Meynet (\\cite{maemey00}). Therefore it becomes of great importance to see whether microturbulence can be responsible for the {\\it He discrepancy}.  In Sect. 2 we present line formation calculations of H and He lines and their behaviour with microturbulence in the O stars domain. In Sect. 3 we determine the effect of including microturbulence in the determination of stellar parameters, and in Sect. 4 we analyze some stars with and without considering microturbulence. Sections 5 and 6 are then dedicated to our discussion and conclusions, respectively . ",
        "conclusions": "We study for the first time the effect of microturbulence in the whole O spectral range, from late to early O types. Introducing microturbulence in the solution of the statistical equilibrium and transfer equations, and then in the formal solution (i.e., in the absorption coefficient), we conclude that for higher gravities the effect in the lines is negligible. This together with the low values of microturbulence usually found for high luminosity class stars, lead us to conclude that only O supergiants have sensitivity to microturbulence effects at a given $T_\\mathrm{eff}$.  In examining the behaviour with microturbulence of the lines we use in our analysis, we show that only \\ion{He}{i} lines and the core of \\ion{He}{ii} $\\lambda$ 4686 \\AA~ are sensitive to microturbulence. For \\ion{He}{i} lines we show that, as should be expected, there is not a constant pattern for each line, but it depends on the parameters considered, that determine the strength of the line and its degree of saturation. This invalidates generalizations to the whole O spectral range made upon results obtained just for a certain spectral type.  In order to quantify the sensitivity of stellar parameters to microturbulence we find that changes in the parameters induced by a value of microturbulence of 15 kms$^{-1}$ are enclosed within the standard error box of our analyses. We think that the lack of a clear pattern in the changes induced in $T_\\mathrm{eff}$ and $\\log g$ is just due to the fact that we are varying the parameters within this error box. We do however find a systematic change in $\\epsilon$ towards lower He abundances when microturbulence is introduced. In particular we find that late O supergiants show a decrease of 0.02--0.04 in $\\epsilon$ (this last value when including other effects that add to microturbulence), which is in agreement with previous results pointing to the inverse relation between the microturbulence assumed and the He abundance obtained (Smith \\& Howarth \\cite{Smith&Howarth98}). Early types are less sensitive to microturbulence, and might not show a difference in the derived He abundance.  Thus microturbulence is not capable of explaining the {\\it He discrepancy} at all for early O stars, and neither it is for late O types with high overabundances.  Looking to individual lines we find that the fits to \\ion{He}{i} $\\lambda$ 4471 \\AA~ are improved when considering microturbulence, but not to the extent of completely explaining its dilution. On the other hand \\ion{He}{i} $\\lambda$$\\lambda$ 4922, 4387 \\AA~ are sometimes slightly better and sometimes slightly worse fitted, in the last case with model cores a bit too strong or a bit too weak, respectively. The rest of the lines keep the same quality in the fit. The {\\emph{\\ion{He}{i} lines problem}} is therefore only partially solved by simply considering microturbulence, even with line--blocking included in the model profiles.  Therefore our conclusion is that microturbulence is affecting the derivation of stellar parameters, but its effect is comparable to the adopted uncertainties. Thus it can reduce moderate He overabundances and solve line fit quality differences, but it cannot explain by itself large He overabundances in O stars, and we are forced to conclude that these are due to other effects, whether real or caused by artifacts in our analyses. This last point will probably not find a definitive answer until we are able to derive reliable abundances of C, N and O in the atmospheres of O stars that we can correlate with the He abundances."
    },
    "0002/astro-ph0002441_arXiv.txt": {
        "abstract": "We study the two-dimensional, time-dependent hydrodynamics of  radiation-driven winds  from luminous accretion disks threaded by a strong, large-scale, ordered magnetic field. The radiation force is mediated primarily by spectral lines and is calculated  using a generalized multidimensional formulation of the Sobolev approximation. The effects of the magnetic field are approximated by adding into the equation of motion a force that emulates a magnetocentrifugal force. Our approach allows us to calculate disk winds when the magnetic field controls the  geometry of the flow, forces  the flow to corotate with the disk, or both. In particular, we calculate models where the  lines of the poloidal component of the field are straight and inclined to the disk at a fixed angle.  Our numerical calculations show that flows which conserve  specific angular velocity have a larger mass loss rate  than their counterparts which conserve specific angular momentum. The difference in the mass loss rate between these two types of winds can be several orders of magnitude for low disk luminosities but vanishes for high  disk luminosities. Winds which conserve angular velocity have  much higher velocities than angular momentum conserving winds. Fixing the wind geometry stabilizes winds which are unsteady when the geometry is derived self-consistently. The inclination angle between the poloidal velocity and the normal to the disk midplane is important. Non-zero inclination angles allow the magnetocentrifugal force to increase the mass loss rate for low luminosities, and increase the wind velocity for all luminosities. The presence of the azimuthal force does not change the mass loss rate when the geometry of the flow is fixed.  Our calculations also show that the radiation force can launch winds from  magnetized disks. The line force  can be essential in producing magnetohydrodynamical (MHD) winds from disks where the thermal energy is too low to launch winds or where the field lines make an angle of $< 30^o$ with respect to the normal to the disk midplane. In the latter case the wind will be less decollimated near its base by the centrifugal force and its collimation far away from the disk can be larger than the collimation of its centrifugally-driven MHD counterpart. ",
        "introduction": "Accretion disks are believed to lose mass via powerful outflows in many astrophysical environments such as active galactic nuclei (AGN); many types of interacting binary stars, non-magnetic cataclysmic variables (nMCVs), for instance; and young stellar objects (YSOs).  Magnetically driven winds from disks are the favored explanation for the outflows in many of these environments. Blandford \\& Payne 1982 (see also Pelletier \\& Pudritz 1992) showed that  the centrifugal force can  drive a wind from  the disk if the poloidal component of the magnetic field, ${\\bf B_p}$ makes an angle of $> 30^o$ with respect to the normal to the disk surface. Generally, centrifugally-driven MHD  disk winds (magnetocentrifugal winds for short) require  the presence of a sufficiently strong, large-scale, ordered magnetic field threading the disk and the poloidal magnetic field to be comparable to the toroidal magnetic field, $|B_\\phi/B_p| \\simless 1$ (e.g., Cannizzo \\& Pudritz 1988, Pelletier \\& Pudritz 1992). Additionally, magnetocentrifugal winds require some thermal assistance  to flow freely and steadily from the surface of the disk,  to pass through a slow magnetosonic surface (e.g., Blandford \\& Payne 1982).  Many authors have  studied  magnetocentrifugal winds from a Keplerian disk (e.g., Ouyed \\& Pudritz 1997, Ustyugova et al. 1999, Krasnopolsky, Li \\& Blandford 1999, and references therein). These studies are either analytic, looking for stationary, often self-similar solutions or numerical, looking for both stationary and time-dependent solutions. However in numerical simulations of magnetocentrifugal winds, the lower boundary is between the slow magnetosonic surface  and the  Alfv${\\acute{\\rm e}}$n surface. This specification requires setting the mass loss rate in advance (e.g., Bogovalov 1997). Simulations of magnetocentrifugal winds  considering the whole disk (even  regions below the slow magnetosonic surface) and not requiring an {\\it ad hoc} mass loss rate, are only now  becoming feasible. This requires, however, an accurate treatment of the radiation and gas pressure effects, among other physical processes.  Thermal expansion and the radiation force have been suggested as other mechanisms that can drive disk winds. These mechanisms can produce powerful winds without the presence of a magnetic field. Winds are likely thermally driven in X-ray-irradiated accretion disks in systems such as  X-ray binaries and AGNs (e.g., Begelman, McKee \\& Shields 1983, Woods et al. 1996). These winds require the gas temperature to be so high that the gas is not gravitationally bound. In such cases, the radiation driving is probably not important because the gas is fully ionized, at least in the hottest regions, and the radiation force is only due to electron scattering. However in the  regions where gas is cooler and not fully ionized the radiation force will be enhanced by spectral lines and  play an important role in controlling the dynamics of the flow. In fact Murray et al. (1995) designed a line-driven disk wind model specifically for AGNs.  Radiation-driven disk winds have been extensively modeled (e.g., Pereyra, Kallman \\& Blondin 1997; Proga, Stone \\& Drew 1998, hereafter PSD~I; Proga 1999; Proga, Stone \\& Drew 1999, hereafter PSD~II). These recent studies showed that radiation pressure due to spectral lines can drive winds from luminous disks. This result has been expected (e.g., Vitello \\& Shlosman 1988). These studies, in particular those by PSD~I, also showed some unexpected results. For example, the flow is unsteady in cases where the disk luminosity dominates the driving radiation field. Despite the complex structure of the unsteady disk wind, the time-averaged mass loss rate and terminal velocity scale with luminosity, as do steady flows obtained where the radiation is dominated by the central object.  To calculate the line force for disk winds, PSD~I adopted the method introduced by Castor, Abbott \\& Klein (1975, hereafter CAK) and further developed by Friend \\& Abbott (1986), Pauldrach, Puls \\& Kudritzki (1986) for one-dimensional radial flows within the context of a  wind from hot, luminous OB stars.  To extend the CAK method to model multi-dimensional disk winds, it is necessary to accommodate the effects of the three-dimensional velocity field and the direction-dependent intensity. Owocki, Cranmer \\& Gayley (1996) showed that these effects can lead to qualitatively different results compared to those obtained from a one-dimensional treatment in the case of a rapidly rotating star. The most difficult aspect of calculating the line force due to a disk is in the evaluation of  the integral involving the velocity gradient tensor over the entire solid angle occupied by the radiating surface. PSD~I used an angle-adaptive quadrature to ensure an accurate result. However, computational limitations required that they simplified the integrand, retaining only the dominant terms in the velocity gradient tensor.  PSD~II generalized the PSD~I method and introduced a new quadrature that avoids any simplification of the integrand.  This allowed them to evaluate the radiation force for completely arbitrary velocity fields within the context of the CAK formalism.   They applied this method to recalculate several disk wind models first discussed in PSD~I. These more physically accurate models show that PSD~I's more approximate method was very robust. The PSD~II calculations predict total mass loss rates and velocities marginally different from those published in PSD~I. Additionally, PSD~II find that models which display unsteady  behavior in PSD~I are also unsteady with the new method. The largest change caused by the new method is in the disk-wind opening angle: winds driven only by the disk radiation are more polar with the new method while winds driven by the disk and central object radiation are typically more equatorial.   In this paper, we will numerically  check how strong, ordered magnetic fields can change disk winds driven by the line force for a given disk luminosity. We assume that the transport of angular momentum in the disk is dominated by local disk viscosity, for instance due to the local shear instability in weakly magnetized disk (Balbus \\& Hawley 1997). Instead of adding the magnetic fields to the PSD~II model and solving consistently the  equations of MHD, we simply start by   adding some of the effects due to the magnetic fields, namely (1) the azimuthal force so the wind  conserves the specific angular velocity along the streamlines and (2) the force perpendicular to the field lines so the geometry of the disk wind is controlled by the magnetic field. We thus adopt the popular concept that the magnetic field dominates outside the disk and at  least near the disk surface, one can think of the magnetic field lines as rigid wires that control any flow (e.g., Blandford \\& Payne 1982; Pelletier \\& Pudritz 1992). Such an approach is clearly simplistic and can be applied only to sub-Alfv${\\acute{\\rm e}}$nic flows. Nevertheless, the results presented here provide a useful exploratory study of  line-driven disk winds in the presence of  a strong, large-scale field magnetic threading the disk.  The organization of the paper is as follows. In Section 2 we describe our numerical calculations; in Section 3 we present our results; in Section 4 we conclude with a brief discussion. ",
        "conclusions": "We have studied winds from accretion disks with very strong organized magnetic fields and the radiation force due to lines. We use numerical methods to solve the two-dimensional, time-dependent equations of hydrodynamics. We have accounted for the radiation force using a generalized multidimensional formulation of the Sobolev approximation. To include the effects of the strong magnetic field we have added to the hydrodynamic equation an approximate magnetocentrifugal force. Our approach of treating the latter allows us to study cases where the wind geometry is controlled entirely by the magnetic field with the lines of the poloidal component of the field, ${\\bf B}_p$ being  straight and inclined at a fixed angle to the disk midplane where the specific angular velocity is conserved along the streamlines due to the magnetic azimuthal force, or both.  Our simulations of line-driven disk winds show that inclusion of the azimuthal force, ${\\bf F}^t$ which makes the wind  corotate with the disk, increases the wind mass loss rate significantly only for low disk luminosities. As expected such winds have much higher velocities than winds with zero azimuthal force. Fixing the wind geometry effectively reduces the flow from two-dimensional to one-dimensional. This in turn stabilizes  winds   which are unsteady when the geometry is allowed to be determined self-consistently. The inclination angle between the poloidal velocity and the normal to the disk midplane is important. If it is higher than zero it can significantly increase the mass loss rate for low luminosities, and increase the wind velocity for all luminosities. The presence of the azimuthal force does not change the mass loss rate if the geometry is fixed.  It is very intriguing that the mass loss rate can be enhanced by the magnetocentrifugal force the most for low luminosities where line force is  weak and at the same time the mass loss rate remains a strong function of the disk luminosity for all luminosities. In particular, there is no wind when $\\dot{M}_a \\simless 10^{-8}~\\MSUNYR$ as in purely line-driven case (e.g.,  PSD~I; Proga 1999).  Model IIa' with $i=0$ is steady and has the same $\\dot{M_w}$ as the corresponding PSD~II model. The mass loss rate is then the same regardless of the time behavior. This is consistent with PSD~I conclusion that the mass loss rate depends predominately on the total system luminosity.  Drew \\& Proga (1999) applied results from PSD~I, PSD~II and Proga~1999 to nMCV. In particular, they compared mass loss rates predicted by the  models with observational constraints. They concluded that either mass accretion rates in high-state nMCV are higher than presently thought  by a factor of 2-3  or that radiation pressure alone is not quite sufficient to drive the observed hypersonic flows. If the latter were true then an obvious candidate to assist radiation pressure in these cases is MHD (e.g., Drew \\& Proga 1999).  An increases of the mass loss rate due to inclusion of a magnetocentrifugal force (our case I) brings our predictions close to the observational estimates for nMCVs. However we should bear in mind that at the same time, the radial velocity of the wind increases above the observed velocities in nMCVs that are of the order a few thousand ${\\rm km~s^{-1}}$. The increase of the wind velocity may not be  large if the magnetic field is moderately strong and spins up the disk wind only close to its base where the mass loss rate is determined. Then farther away from the disk the wind would  gain less or no angular momentum and velocities would be similar to those in the case without the magnetic field at all.  Outflows which conserve  specific angular velocity have a higher rotational velocity than those which conserve  specific angular momentum. The former also have  higher terminal velocities due to stronger centrifugal force. Additionally in the wind where the specific angular velocity is conserved the rotational velocity is comparable to the terminal velocity while in the wind where the specific angular momentum is conserved the rotational velocity decreases asymptotically to zero with increasing radius and is therefore much lower than the terminal velocity. Thus we should be able to distinguish these two kinds of winds based on their line profiles. For example, highly rotating  winds should produce emission lines  much broader than slowly rotating, expanding winds, if we see the disk edge-on.  So far we have discussed changes of the line-driven wind caused by inclusion of the magnetocentrifugal force. However we can also anticipate some changes in the magnetocentrifugal wind if we add the line force. For example, there may be  a difference in the wind collimation caused by the line force. The centrifugal force decollimates the magnetocentrifugal outflow near the disk as $i$ must be greater than $30^o$. In such  a case the magnetic field collimates the outlow only beyond the Alfv${\\acute{\\rm e}}$n surface where the pinch  force exerted by the toroidal component of the field operates. In the case with a significant line driving, the magnetic field does not have to be inclined at the angle $> 30^o$. Then in such case, outflows are more collimated than pure magnetocentrifugal outflows from the start and may end up more collimated far away from the disk.  Let us now discuss some results of  other related works. Recently Ogilvie \\& Livio (1998) studied a thin accretion disk threaded by a poloidal magnetic field. Their purpose was to determine qualitatively how much thermal assistance is required for the flow to pass through the slow magnetosonic surface.  They found that a certain potential difference must be overcome even when $i > 30^o$ and that thermal energy is not sufficient to launch an outlow from a magnetized disk. Ogilvie \\& Livio suggested that an additional source of energy, such as coronal heating may be required. PSD~I showed the radiative line force can  drive disk winds. Thus the 'missing' energy  may be in the radiation field. Our calculations here show that if the disk luminosity is sufficiently high the line and magnetocentrifugal forces produce strong transonic winds regardless of the wind geometry. Then the magnetocentrifugal force can assist the line force in producing disk winds -- our approach to the problem or the line force or/and the thermal force can well assist  MHD -- Ogilvie \\& Livio's approach (see also Wardle \\& K$\\ddot{\\rm o}$ngil 1993; Cao \\& Spruit 1994, for instance). In seeking a steady state solution of a line-driven disk wind, Vitello \\& Shlosman (1988) also found that the flow must overcome a potential difference -- an increase of the vertical gravity component with height. They suggested that for the radiation force to increase with height above the disk midplane and overcome this potential difference a very particular variation in the ionization state of the gas is required.  There are a number of limitations of our treatment of magnetic fields which are worthy of mention.  Our models  do not include the whole richness of MHD because we approximate the Lorentz force by the 'magnetocentrifugal' force (eqs 3-8) instead of solving self-consistently the equation of motion and the induction equation. We have included effects of a magnetic field on the disk winds in a  simplistic manner that mimics some effects of a very strong, organized magnetic field where $B_\\phi \\simless B_p$ near the disk surface. The magnetic field is treated as a rigid wire that controls the flow  geometry outside the disk. More quantitatively, this corresponds to the situation where outside the disk, at least  near the disk photosphere, the magnetic field dominates i.e., the magnetic pressure  exceeds the disk gas pressure $\\beta\\equiv 8\\pi p_D/ B^2_D < 1 $. Inside the disk however the situation may be different. We assume that the disk is in a steady state that is also stable. In such a case, the regions near the disk midplane are likely supported by the gas pressure and the magnetic pressure should be  smaller than the gas pressure, i.e., $\\beta   \\simgreat 1$, otherwise the disk may be unstable (Stella \\& Rosner 1984). Using the system parameters adopted here: $c_s= 14~\\rm km~s^{-1}$ and $\\rho_o=10^{-4}~\\rm g~cm^{-3}$, we find that the $\\beta \\simgreat 1$ condition yields a maximum value for the disk magnetic fields strentgh of the order of $\\sim~10^5$~G. Our assumption that poloidal magnetic field lines are straight excludes any magnetic tension in the $(r,\\theta)$ plane.  In our cases~I and III we assume that the specific angular velocity, $\\Omega$ is conserved  along a streamline from the footpoint of the line to infinity or rather to  the outer boundary of the computational domain. However in a steady state axisymmetric MHD flow, the quantity which is conserved is the total angular momentum per unit mass which can be written as \\begin{equation} l=\\Omega r^2 \\sin^2\\theta - \\frac{r \\sin\\theta B_\\phi} {4 \\pi\\kappa} \\equiv \\Omega_D r_A^2, \\end{equation} where $r_A$ is the position of the the Alfv${\\acute{\\rm e}}$n point of the flow on the streamline, and $\\kappa=\\rho v_p/B_p$ is the mass load-- another constant along the streamline (e.g., Pelletier \\& Pudritz 1992). The total angular momentum has contributions from both the flowing rotating gas and the twisted magnetic field. This means that as the material angular momentum, $\\Omega r^2 \\sin^2\\theta $ increases due to corotation with the disk, the toroidal component of the magnetic field must increase so the total angular momentum is conserved. Then our approximation in cases~I and III corresponds to the situation where near the disk the total angular momentum is dominated by the contribution from the twisted magnetic field and  the Alfv${\\acute{\\rm e}}$nic surface is beyond our computational domain. In other words our approximation is valid in the region where the outflow is sub-Alfv${\\acute{\\rm e}}$nic. Consequently,  our approach  does not allow us to study collimation of the wind by the magnetic field because this happens beyond the fast magnetosonic surface where the flow is super-Alfv${\\acute{\\rm e}}$nic.  PSD~II's models  and ours in case II correspond to the other extreme case where the total angular momentum has no contribution from the twisted magnetic field and the material angular momentum is conserved.  Our treatment of magnetic fields does not include effects of the magnetic pressure. The magnetocentrifugal driving presupposes the existence of a strong poloidal magnetic field comparable to the toroidal magnetic field near the disk surface,  $|B_\\phi/B_p| \\simless 1$.  However when the poloidal magnetic field is weaker than the toroidal magnetic field, $|B_\\phi/B_p| >> 1$ the magnetic pressure may be dynamically important in driving disk winds (e.g., Uchida \\& Shibata 1985; Pudritz \\& Norman 1986, Shibata \\& Uchida 1986, Contopoulos 1995; Kudoh \\& Shibata 1997; Ouyed \\& Pudritz 1997 and references therein). For $|B_\\phi/ B_p| >> 1$, there is initially the buildup of the toroidal magnetic field by the differential rotation of the disk that in turn generates the magnetic pressure of the toroidal field. The magnetic pressure then gives rise to a self-starting wind. To produce a steady outflow driven by the magnetic pressure a steady supply of the advected toroidal magnetic flux at the wind base is needed, otherwise the outflow is likely a transient (e.g., K$\\ddot{\\rm o}$nigl 1993, Contopoulos 1995, Ouyed \\& Pudritz 1997). However it is not clear whether the differential rotation of the disk can produce such a supply of the toroidal magnetic flux to match the escape of magnetic flux in the wind and even if it does whether such a system will be stable (e.g., Contopoulos 1995, Ouyed \\& Pudritz 1997 and references therein).   Concluding, we would like to stress that a further development of models of radiation-driven winds from disks should include taking into account magnetic field but it is equally important to consider adding the radiation force to models of MHD winds from luminous disks. Both kinds of models are quite well understood now and if merged they could allow us to study better disk winds in systems such as cataclysmic variables and AGNs. As we mentioned above thermal assistance may not be sufficient to launch wind from a magnetized disk and we should consider not only coronal heating but also line driving.  ACKNOWLEDGEMENTS: We thank John Cannizzo, Janet Drew, Achim Feldmeier, Tim Kallman, Scott Kenyon, Mario Livio, and James Stone for comments on earlier drafts of this paper. We also thank Steven Shore and an anonymous referee for comments that helped us clarify our presentation. This work was performed while the author held a National Research Council Research Associateship at NASA/GSFC. Computations were supported by NASA grant NRA-97-12-055-154.  \\newpage"
    },
    "0002/astro-ph0002394_arXiv.txt": {
        "abstract": "Exact calculations of innermost stable circular orbit (ISCO) around rotating strange stars are performed within the framework of general relativity. Equations of state (EOS) of strange quark matter based on the MIT Bag Model with massive strange quarks and lowest order QCD interactions, are used. The presence of a solid crust of normal matter on rotating, mass accreting strange stars  in LMXBs  is taken into account. It is found that, contrary to neutron stars, above some minimum mass (which for the considered equations of state ranged from $1.4~{\\rm M}_\\odot$ to $1.6~{\\rm M_\\odot}$) a gap always separates the ISCO and stellar surface, independently of the strange star rotation rate. For a given baryon mass of strange star, we calculate the ISCO frequency as function of stellar rotation frequency, from static to Keplerian configuration. For masses close to the maximum mass of static configurations the ISCO frequencies  for static and  Keplerian configurations are similar. However, for masses significantly lower than  the maximum mass of static configurations, the minimum value of the ISCO frequency is reached in the Keplerian limit. Presence of a solid crust increases the ISCO frequency for the Keplerian configuration by about  ten percent compared to that for a bare strange star of the same mass. For standard parameters of  strange quark matter EOS, resulting in a maximum static mass of $1.8~{\\rm M}_\\odot$, the ISCO frequency  for $\\ga 1.4~{\\rm M}_\\odot$ strange stars always exceeds 1.07~kHz (the upper QPO frequency  reported for 4U 1820-30). We give an example of strange quark matter model, which yields maximum static mass of $2.3~{\\rm M}_\\odot$, and for which the ISCO frequency of 1.07 kHz is allowed at stellar rotation rates 200-300 Hz, provided the strange star mass exceeds 2.2 ${\\rm M}_\\odot$. For this EOS even lower  value $\\nu_{\\rm ISCO}\\simeq 1$~kHz is reached near  the Keplerian limit, for a broad range of stellar masses. While reproducing $\\nu_{\\rm ISCO}=1.07$ kHz at slow rotation rates requires  tuning of strange quark matter parameters, no such a tuning is required to reproduce orbital frequencies around strange stars equal to highest observed upper QPO frequencies. ",
        "introduction": "Observations of quasi periodic oscillations (QPOs) in the X-ray fluxes from low-mass X-ray binaries (LMXB), which are believed to be due to the orbital motion of matter in an accretion disk, raised hopes concerning observational constraints on the equation of state (EOS) of matter at supranuclear densities (Kaaret et al. 1997, Klu{\\'z}niak 1998, Zhang et al. 1998, Miller et al. 1998, Thampan et al. 1999, Schaab \\& Weigel 1999). General relativity predicts the existence of the marginally stable (MS) orbit, within which no stable circular motion is possible. This implies the existence of the innermost stable circular orbit (ISCO) around neutron stars. The frequency of the ISCO is an upper bound on the frequency of stable orbital motion around neutron stars. Whether the ISCO is separated from neutron star surface by a gap, or its radius coincides with stellar equatorial radius, depends on the star mass and on the EOS of neutron star matter. On the other hand, accreting neutron stars in LMXBs are expected to be rotating, and this influences both neutron star structure and the ISCO. Therefore, in order to attempt to use observed frequencies of QPOs to constrain the EOS of dense matter, one has to calculate the ISCO as a function of stellar mass and stellar rotation frequency. Such a procedure is based on the assumption that the observed upper QPO frequency is due to orbital motion, and that the effects of magnetic field, accretion, and radiation drag on the matter flow can be neglected.  A basic assumption of the present paper is that the frequency of the upper kHz QPO is the orbital frequency of the inner edge of an accretion disk  surrounding the compact object, which will be identified with the ISCO. This is the leading interpretation of the QPOs. However, alternative models of the kHz QPOs were also proposed. In a model of Alpar \\& Yilmaz (1997) the kHz QPOs are explained in terms of wave packets of sound waves  in the inner disk. In a series of papers, Titarchuk and collaborators propose a model in which the QPOs result from radial oscillations of the plasma in the boundary layer, i.e. in the region between the ISCO and stellar surface (see Titarchuk \\& Osherovich 1999 and references therein). These alternative models will not be considered in our study.  \\begin{figure*}     % \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{zdunik1.eps}} \\caption{ The radius of the ISCO (dotted line), equatorial radius of a bare strange star (solid line), equatorial radius of a strange star with maximally thick solid crust (long-dashed line), versus rotation frequency of strange star. Thin short-dashed  line corresponds to slow-rotation approximation of Klu{\\'z}niak \\& Wagoner (1985). Calculations  were performed for the SQM1 EOS of strange matter (see the text). Figures  correspond to stellar models with  fixed total baryon number, equal to that of a static star of  gravitational mass of $1.2,~1.4,~1.6, 1.75~{\\rm M}_\\odot$. Maximum mass for static strange stars is $1.8~{\\rm M}_\\odot$ } \\label{fig1} \\end{figure*}  In the present paper we describe  results of exact  calculations of the ISCO, under the assumption that the compact object is not a neutron star, but a strange star. A strange star is composed of self-bound quark matter, which at zero pressure would constitute a real ground state of matter (strange matter), with energy per unit baryon number lower  than that of $^{56}{\\rm Fe}$ crystal (Witten 1984, Farhi \\& Jaffe 1984, Haensel et al. 1986, Alcock et al. 1986; for a recent review of physics and astrophysics of strange matter, see Madsen 1999). Recently, strange stars were invoked by several authors in the context of modeling of observational properties of some X-ray and gamma-ray sources (Bombaci 1997, Cheng et al. 1998, Dai \\& Lu  1998, Li et al. 1999). First  study of possibility of existence of strange stars  in LMXBs  exhibiting  kHz QPOs was restricted to slow-rotation approximation for the ISCO, neglected the effect of rotation on the strange star  structure, and used simplified EOS of strange matter, with massless, non-interacting quarks (Bulik et al. 1999 ). Very recently, the ISCOs around bare strange stars were calculated, assuming a simplified EOS of strange matter,  for the limiting case of rotation at Keplerian frequency (Stergioulas et al. 1999). In both these studies the possible presence of the solid crust on the strange star  surface was not taken into account.  In principle, a strange star  could be covered by a thin   crust of normal matter, a possibility which is particularly natural in the case of LMXB. The problem of formation and structure of a crust on an accreting  strange star  was studied by  Haensel \\& Zdunik (1991)(see also Miralda-Escud{\\'e} et al. 1990). Because of its low mass, typically $\\la 10^{-5}~{\\rm M}_\\odot$, the effect of the crust on the exterior spacetime is negligible. However,  it determines the location of the star surface, due to its finite thickness of $\\sim 200-300~$m. The matter distribution within the strange core, relevant for the exterior metric of rotating strange star, is characterized by a very flat density profile: for a massive strange star, density at the stellar center is typically only 2-3 times larger than that at the outer edge of the strange core. This has to be contrasted with the density distribution within a massive neutron star, which decreases continuously from $\\sim 10^{15}~{\\rm g~cm^{-3}}$ at the center to a few ${\\rm g~cm^{-3}}$ at the surface. The differences between the density profiles of a strange star and a neutron star result from the basic difference in the EOS of their interiors. In the case of a rapidly rotating compact object (situation relevant to LMXB), the differences between matter distributions within neutron star and strange star may be  expected to  imply differences in the spacetime exterior to the compact object, and in particular, differences in the properties of the ISCO. ",
        "conclusions": "The features of ISCOs around rapidly rotating strange stars, described in the present paper for a particular choice of strange matter EOS, are actually generic. The MIT Bag Model EOS of strange quark matter depends on $B$, $m_{\\rm s}$, and $\\alpha_{\\rm c}$ in a way, which implies specific scaling properties with respect to change of $B$ (Haensel et al. 1986, Zdunik \\& Haensel 1990). As a consequence, the global parameters of rotating strange stars  scale with some power of $B$, which allows one to determine the values of $M$, $R_{\\rm eq}$, $R_{\\rm ms}$, etc., for $B$, $\\alpha_{\\rm c}$, and $m_{\\rm s}$, from those calculated for $B_0$, $\\alpha_{\\rm c}$ and strange quark mass $m_{\\rm s}(B_0/B)^{1/4}$. All  length-type quantities (stellar radius, thickness of the crust and radius of the ISCO)  scale as $B^{-1/2}$, and all frequencies ($\\nu_{\\rm ISCO}$, $\\nu_{\\rm rot}$) scale as $B^{1/2}$, e.g., $\\nu_{\\rm ISCO}[B]=\\nu_{\\rm ISCO}[B_0]\\cdot (B/B_0)^{1/2}$ and $R_{\\rm ms}[B]=R_{\\rm ms}[B_0]\\cdot (B/B_0)^{-1/2}$. Thus, for other values of $B$ the patterns of lines in Figs. 1-4 do not change, provided one rescales the axes and stellar masses.  Our calculations show that the properties of the ISCOs around strange stars differ from those around neutron stars. A generic property is the existence of the gap between the ISCO and the stellar surface, for both slowly and rapidly rotating strange stars.  The highest observed QPO frequency of {\\chng 1.33 kHz in 4U 0614+91} can be easily interpreted as an orbital frequency around strange star based on the standard SQM1 EOS of strange matter, with no significant constraint on strange star mass and rotation rate. In the case of the SQM2 EOS, the orbital origin of the {\\chng 1.33 kHz QPO implies $M\\la 1.8~{\\rm M}_\\odot$ at rotation frequencies $\\sim 300$~Hz}, while frequencies close to the mass shedding limit are excluded.  The value of the ISCO frequency at the Keplerian limit is significantly influenced by the presence of a crust  on the strange star surface, which increases this frequency by about  hundred Hz compared to the value for a bare strange star of the same mass. As one expects the presence of a crust on a strange star in a LMXB, we conclude that only slowly rotating strange stars with  mass above $2.2~{\\rm M}_\\odot$ seem to be consistent with $\\nu_{\\rm ISCO}=1.07$~kHz. This excludes EOS of strange matter corresponding to the standard bag model parameters, and can be satisfied only by choosing a set of parameters quite different from the standard one.   The numerical results discussed in the present paper show that consistency of the ISCOs around slowly rotating strange stars with orbital-motion interpretation of QPOs in LMXBs can be achieved only with a substantial tuning of the MIT Bag Model parameters of strange matter. Our SQM2 EOS is a result of such a tuning. For this EOS, the condition $\\nu_{\\rm ISCO}\\simeq 1$~kHz is satisfied not only for slowly rotating massive models with $M\\ga 2.2~{\\rm M}_\\odot$, but also for a broad range of masses of configurations close to the Keplerian  limit. In contrast to $\\nu_{\\rm ISCO}$ at low rotation rate, which decreases with increasing baryon mass, the ISCO frequency at the Keplerian limit decreases with decreasing baryon mass of rotating strange star."
    },
    "0002/astro-ph0002284_arXiv.txt": {
        "abstract": "A Cosmic Far-InfraRed Background (CFIRB) has long been predicted that would traces the intial phases of galaxy formation. It has been first detected by~\\cite{PUG96} using COBE data and has been later confirmed by several recent studies (~\\cite{FIX98},~\\cite{HAU98},~\\cite{LAG99}). We will present a new determination of the CFIRB that uses for the first time, in addition to COBE data, two independent gas tracers: the HI survey of Leiden/Dwingeloo (~\\cite{HAR98}) and the WHAM H$_{\\alpha}$ survey (~\\cite{REY99}). We will see that the CFIRB above 100 $\\mu$m is now very well constrained. The next step is to see if we can detect its fluctuations. To search for the CFIRB fluctuations, we have used the FIRBACK observations. FIRBACK is a deep cosmological survey conducted at 170$\\mu$m with ISOPHOT (~\\cite{DOL00}). We show that the emission of unresolved extra-galactic sources clearly dominates, at arcminute scales, the background fluctuations in the lowest galactic emission regions. This is the first detection of the CFIRB fluctuations. ",
        "introduction": "In very diffuse parts of the sky (no molecular clouds and HII regions), the far-IR emission can be written as the sum of dust emission associated with the neutral gas, dust associated with the ionised gas, interplanetary dust emission and the CFIRB (and eventually the cosmological dipole and CMB). In previous studies (~\\cite{PUG96},~\\cite{FIX98},~\\cite{HAU98}), dust emission associated with the ionised gas which was totally unknown has been either not subtracted properly or neglected. \\\\ We have detected for the first time dust emission in the ionised gas (~\\cite{LAG99}) and shown that the emissivity (which is the IR emission normalised to unit hydrogen column density) of dust in the ionised gas was nearly the same as that in the neutral gas. This has consequences on the determination of the CFIRB. Following this first detection, we have combined HI and WHAM H$_{\\alpha}$ data (~\\cite{REY99}) with far-IR COBE data in order to derive dust properties in the diffuse ionised gas as well as to make a proper determination of the CFIRB. Technically, after a careful pixel selection (see~\\cite{LAG00} for more details) we describe the far-infrared dust emission as a function of the HI and H$^+$ column density by: \\begin{equation} \\label{main_EQ} IR= A \\times N(HI)_{20cm^{-2}} + B \\times N(H^+)_{20cm^{-2}} + C \\end{equation} where N(HI)$_{20cm^{-2}}$ and N(H$^+$)$_{20cm^{-2}}$ are the column densities normalised to 10$^{20}$ H cm$^{-2}$. The coefficients A, B and the constant term C are determined simultaneously using regression fits. We show that about 25$\\%$ of the IR emission comes from dust associated with the ionised gas which is in very good agreement with the first determination (~\\cite{LAG99}). The CFIRB spectrum obtained using this far-infrared emission decomposition is shown in Fig.~\\ref{CFIRB_spec} together with the CFIRB FIRAS determination of~\\cite{LAG99} in the Lockman Hole region. We see a very good agreement between the two spectra. These determinations are also in good agreement with~\\cite{FIX98}.  \\begin{figure} \\centering \\includegraphics[width=.8\\textwidth]{glagache_f1.ps} \\caption[]{CFIRB spectra obtained from the decomposition of the far-infrared sky (continuous line) and determined for the Lockman Hole region (dashed line) by~\\cite{LAG99}. Also reported are DIRBE values at 100, 140 and 240 $\\mu$m.} \\label{CFIRB_spec} \\end{figure}  At 140 and 240 $\\mu$m, the values obtained for the CFIRB are 1.13$\\pm$0.54 MJy/sr and 0.88$\\pm$0.55 MJy/sr respectively For each selected pixel, we compute the residual emission, R = IR - A$\\times$N(HI) - B$\\times$N(H$^+$). Uncertainties of the CFIRB have been derived from the width of the histogram of R (statistical uncertainties derived from the regression analysis are negligible). The obtained CFIRB values, although much more noisy (due to the small fraction of the sky used), are in very good agreement with the determination of ~\\cite{HAU98}. At 140~$\\mu$m, the CFIRB value of~\\cite{LAG99} is smaller than that derived here since the assumed WIM (Warm Ionised gas) dust spectrum was overestimated (the WIM dust spectrum was very noisy below 200~$\\mu$m and the estimated dust temperature was too high).\\\\  At 100~$\\mu$m, assuming an accurate subtraction of the zodiacal emission, our decomposition gives: I$_{CFIRB}$(100)= 0.78$\\pm$0.21 MJy/sr. This is the first time that two independent gas tracers for the HI and the H$^+$ have been used to determine the background at 100 $\\mu$m. One has to note that methods based on the intercept of the far-IR/HI correlation for the determination of the CFIRB are dangerous. For example, for our selected parts of the sky, this intercept is about 0.91 MJy/sr, which is quite different from the value of the CFIRB (0.78 MJy/sr). The CFIRB value of 0.78 MJy/sr can be compared to the non-isotropic residual emission found by~\\cite{HAU98}. The average over three regions of the residual emission, equal to 0.73$\\pm$0.20 MJy/sr, is in very good agreement with our determination.\\\\  So we see, using different approaches, that we are now converging on the shape and level of the CFIRB above 100 $\\mu$m. The next step is to see if we can detect its fluctuations and study them. ",
        "conclusions": "We have shown in the FIRBACK fields that the extra-galactic background fluctuations lie well above the instrumental noise and the cirrus confusion noise. The observed power spectrum shows a flattening at high spatial frequencies which is due to unresolved extra-galactic sources. The level of the extra-galactic power spectrum fluctuations is nearly the same in all FIRBACK fields. The next step consists of removing the cirrus contribution using independent gas tracers (the H$_{\\alpha}$ and the 21cm emission lines) to isolate the extra-galactic fluctuation brightness and try to constrain the IR large scale structures.\\\\"
    },
    "0002/hep-ph0002064_arXiv.txt": {
        "abstract": "Results of neutrino oscillation experiments have always been presented on the $(\\sin^{2} 2\\theta, \\Delta m^{2})$ parameter space for the case of two-flavor oscillations.  We point out, however, that this parameterization misses the half of the parameter space $\\frac{\\pi}{4} < \\theta \\leq \\frac{\\pi}{2}$ (``the dark side''), which is physically inequivalent to the region $0 \\leq \\theta \\leq \\frac{\\pi}{4}$ (``the light side'') in the presence of matter effects.  The MSW solutions to the solar neutrino problem can extend to the dark side, especially if we take the conservative attitude to allow higher confidence levels, ignore some of the experimental results in the fits, or relax theoretical predictions.  Furthermore, even the so-called ``vacuum oscillation'' solution distinguishes the dark and the light sides.  We urge experimental collaborations to present their results on the entire parameter space. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002147_arXiv.txt": {
        "abstract": "We present accurate UBV(RI)$_{\\rm C}$ photometric sequences around 20 symbiotic stars. The sequences extend over wide brightness and color ranges, and are suited to cover quiescence as well as outburst phases. The sequences are intended to assist both present time photometry as well as measurement of photographic plates from historical archives. \\keywords {Catalogs -- Binaries: symbiotic} ",
        "introduction": "Symbiotic stars are binary systems composed of a cool giant and a hot, luminous white dwarf. They show variability over any time scale from minutes (flickering) to several decades (outbursts of symbiotic novae), with phenomena related to the orbital motion having periodicities generally between 1 and 4 years (or a few decades in the systems harbouring a Mira variable, $\\sim$20\\% of all known symbiotic stars).  Such long time scales tend to discourage stand--alone photometric campaigns from a single Observatory (which would require observational programs running up to $\\sim$10 years in order to derive - for example - a firm orbital period). Most of the current photometric investigations of symbiotic stars therefore try to assemble as much as possible data from the widest set of current and archival sources. A template example is the recent reconstruction of the 1890-1996 lightcurve and orbital period determination for YY~Her by Munari et al. (1997).  The data so collected are generally very heterogeneous in nature, with large differences caused for example by ($a$) the non-standard photometric bands, ($b$) the adopted comparison sequences and standard stars, ($c$) lack of adherence to and transformation into a system of general use, like the UBV(RI)$_{\\rm C}$, and ($d$) telescope focal length or pixel scale that causes blending with images of nearby field stars. These differences generally may introduce such a large scatter in the data that all but the strongest details are washed out.  The establishment of suitable and accurate photometric comparison sequences covering a wide range in magnitude and colors should alleviate considerably some of the above problems, and could encourage small observatories and/or occasional observers to obtain new data as well as to encourage those with access to plate archives to search for valuable historical data.  To this aim we present here suitable, UBV(RI)$_{\\rm C}$ comparison sequences for 20 symbiotic stars (all but a few accessible from both hemispheres. See Table~1 for a list of the program stars). The sequences are basically intended to allow a general observer to capture on a single CCD frame or to have in the same eyepiece field of view when inspecting archival photographic plates: ($a$) enough stars to cover the whole range of known or expected variability for the given symbiotic star, ($b$) stars of enough different colors to be able to calibrate the instrumental color equations and therefore reduce to the standard UBV(RI)$_{\\rm C}$ system the collected data, and ($c$) stars well separated from surrounding ones to avoid blending at all but the shortest telescope focal lengths.  \\begin{table*} \\caption[]{List of program symbiotic stars. The coordinates for the symbiotic stars are from our observations (equinox J2000.0, mean epoch 1999.5). The $e_\\alpha$ and $e_\\delta$ columns list the errors in milliarcsec for right ascension and declination, respectively. The last two columns list the coordinates of the field centers in Figures~1 and 2.} \\centerline{\\psfig{file=tab_1.ps,width=18cm}} \\end{table*}  \\begin{figure*}[h!] \\centerline{\\psfig{file=fig_1.ps,width=17.5cm}} \\caption[]{Finding charts for the UBV(RI)$_{\\rm C}$ comparison photometric sequence around the program symbiotic stars. The fields are in the same order as in Table~1. North is up and East to the left, with an imaged field of view of 5.16$\\times$5.16 arcmin and a 5.4$\\times$5.4 coordinate grid (see bottom-right panel of Figure~2). Stars are plotted as open circles of diameter proportional to the brightness in the $V$ band. The stars making up the photometric sequence (see Table~2) are plotted as filled circles.} \\end{figure*} ",
        "conclusions": ""
    },
    "0002/astro-ph0002371_arXiv.txt": {
        "abstract": "We study the effect of a single, instantaneous starburst in a gas-rich dwarf galaxy on the dynamical and chemical evolution of its interstellar medium. We consider the energetic input and the chemical yields originating from SNeII, SNeIa and intermediate-mass stars. We find that a galaxy resembling IZw18 develops a galactic wind carrying out mostly the metal-rich gas. The various metals are lost differentially and the metals produced by the SNeIa are lost more efficiently than the others. As a consequence, we find larger [$\\alpha$/Fe] ratios for the gas inside the galaxy than for the gas leaving the galaxy. Finally we find that a single burst occurring in primordial gas (without pre-enrichment), gives chemical abundances and dynamical structures in good agreement with what observed in IZw18 after $\\sim$ 29 Myr from the beginning of star formation. ",
        "introduction": "Blue compact dwarf galaxies (BCD) are gas-rich systems experiencing an intense star formation. These galaxies have very simple structures, small sizes and are very metal poor. For these reasons, BCD are excellent laboratories to investigate the effect of a starburst on the chemical and dynamical evolution of the interstellar medium (ISM).  Previous dynamical and chemical studies of these galaxies have suggested the existence of a `differential galactic wind', in the sense that after a starburst event these objects would loose mostly metals (ref. Mac Low \\& Ferrara 1999; D'Ercole \\& Brighenti 1999; Pilyugin 1992, 1993; Marconi et al. 1994). However, in none of these studies, detailed chemical and dynamical evolution was taken into account at the same time. The aim of this paper is to include the effects (both energetic and chemical) of type II and type Ia SNe in a detailed dynamical model. ",
        "conclusions": ""
    },
    "0002/astro-ph0002237_arXiv.txt": {
        "abstract": "High-resolution optical spectra of the R Coronae Borealis (RCB) star V854\\,Centauri in the early stages of a decline show, in addition to the features reported for other  RCBs in decline, narrow absorption lines from the C$_2$ Phillips system. The low rotational temperature, T$_{\\rm rot}$ = 1150K, of the C$_2$ ground electronic state suggests the cold gas is associated with the developing shroud of carbon dust. These absorption lines were not seen at a fainter magnitude on the rise from minimum light nor at maximum light. This is the first detection of cold gas around a RCB star. ",
        "introduction": "R Coronae Borealis stars are H-poor F-G type supergiants that decline in brightness unpredictably  by up to  8 magnitudes and remain  below their normal brightness for several weeks to months. It is generally accepted that these declines are due to formation of a cloud of carbon soot that obscures the stellar photosphere. Unanswered questions  remain: `What triggers cloud formation?' and `Where does the soot form?' High-resolution spectroscopic monitoring of RCBs from maximum light into decline will likely be necessary  to refine schematic ideas into answers  that are accorded widespread acceptance. We report the first detection of cool gas (T $\\simeq 1100$K) during the early decline of a RCB star and, hence, evidence for a site of soot formation. Cold dust is, of course, known around  RCBs through  detection of an infrared excess.  The RCB in question is V854\\,Cen, which at maximum light is the third brightest RCB variable after R\\,CrB and RY\\,Sgr. V854\\,Cen is presently  the most variable of all Galactic  RCBs. Despite the combination of favorable apparent magnitude and propensity to fade, there is a dearth of high-quality spectroscopic observations of this star in decline. The sole report of a high-resolution optical spectrum  covering a broad bandpass in a deep decline is that by Rao \\& Lambert (1993) taken when the star had faded by about 8 mag. Low resolution spectra are described by Kilkenny \\& Marang (1989) and spectropolarimetric observations are discussed by Whitney et al. (1992). Spectra at high-resolution at maximum light were used by Asplund et al (1998) for their abundance analysis that confirmed that V854\\,Cen has a somewhat unusual composition among the RCBs for which abundance anomalies are a {\\it sine qua non}. In particular, V854\\,Cen, although hydrogen-poor relative to normal stars, is the most hydrogen-rich RCB by a clear margin. Despite limited temporal coverage, our new  spectra of V854\\,Cen in decline provide a  novel result  - the detection of cold C$_2$ gas. Our  spectra  otherwise closely resemble those of the RCBs extensively studied in decline: R\\,CrB (Rao et al. 1999) and RY\\,Sgr (Alexander et al. 1972). This concordance, which  suggests  that RCBs  have a common general structure of their upper atmospheres and circumstellar regions, is  briefly demonstrated here but we focus on the novel  lines of the C$_2$ molecule. ",
        "conclusions": "For the first time, cold gas below the temperature required for soot formation has been detected around a RCB in its decline to a deep minimum. Our detection of absorption by cold C$_2$ molecules around V854\\,Cen now needs to be followed by synoptic observations of this active RCB variable in order to trace the evolution of the cold gas and to place it relative to the star. Is the gas located behind a shock where, as some theories would suppose, dust formation is triggered? Or is it merely an innocent companion to the dust? We recognise that providing synoptic observations at an adequate spectral resolution and high temporal frequency is  a substantial challenge. If the challenge can be met, the result will be a window into the time and place of dust formation, and, perhaps, provie the long-sought understanding of how the characteristic declines of RCBs are initiated."
    },
    "0002/hep-ex0002053_arXiv.txt": {
        "abstract": "{The IGEX Collaboration enriched $^{76}$Ge double-beta decay detectors are currently operating in the Canfranc Underground Laboratory with an overburden of 2450 m.w.e. A recent upgrade has made it possible to use them in a search for WIMPs. A new exclusion plot, $\\sigma$(m), has been derived for WIMP-nucleon spin-independent interaction. To obtain this result, 30 days of data from one IGEX detector, which has an energy threshold $E_{thr} \\sim 4$ keV, have been considered. These data improve the exclusion limits derived from other germanium diode experiments in the $\\sim$50~GeV DAMA region, and show that with a moderate improvement of the background below 10 keV, the DAMA region may be tested with an additional 1~kg-year of exposure.} ",
        "introduction": "Substantial evidence exists suggesting most matter in the universe is dark, and there are compelling reasons to believe it consists mainly of non-baryonic particles. Among these candidates, Weakly Interacting Massive and neutral Particles (WIMPs) are among the front runners. The lightest stable particles of supersymmetric theories, like the neutralino, describe a particular class of WIMPs\\cite{Gri}.  Direct detection techniques rely on measurement of WIMP elastic scattering off target nuclei in a suitable detector\\cite{Mor99}. Slow moving ($\\sim300$ km/s) and heavy ($10 - 10^3$ GeV) galactic halo WIMPs could make a Ge nucleus recoil with a few keV, at a rate which depends on the type of WIMP and interaction. Only about 1/4 of this energy is visible in the detector. Because of the low interaction rate and the small energy deposition, the direct search for particle dark matter through scattering by nuclear targets requires ultralow background detectors with very low energy thresholds.  To detect the possible presence of WIMPs, the predicted event rate is compared with the observed spectrum. If this predicted event rate is larger than the measured one, the particle under consideration can be ruled out as a dark matter component. Such absence of WIMPs can be expressed as a contour line, $\\sigma$(m), on the WIMP-nucleus elastic scattering cross-section plane. This excludes, for each mass, those particles whose cross-section lies above the contour line, $\\sigma$(m).  This direct comparison of the expected signal with the observed background spectrum can only exclude or constrain the cross-section in terms of exclusion plots of $\\sigma$(m). A convincing proof of the detection of dark matter would require finding unique signatures in the data, characteristic of the WIMP, which cannot be attributed to the background or instrumental artifacts. An example is the predicted summer-winter asymmetry \\cite{Dru86} in the WIMP signal rate due to the periodicity of the relative Earth-halo motion resulting from the Earth's rotation around the Sun.  Germanium detectors used for double-beta decay searches have reached one of the lowest background levels of any type of detector and  have a reasonable quenching factor ($\\sim 0.25$). Thus, with sufficiently low energy thresholds, they are attractive devices for dark matter searches.  Germanium diodes dedicated to double-beta decay experiments[4-11] were applied to WIMP searches as early as 1987. The exclusion contour based on the best combination of data from these experiments is referred to in this paper as the ``combined germanium contour''. Only recently has this exclusion plot been surpassed by a sodium iodide experiment\\cite{Ber96} (DAMA NaI-0), which uses a statistical pulse-shape discriminated background spectrum.  This paper presents a new germanium detector data limit for the direct detection of non-baryonic particle dark matter in the $\\sim$50~GeV DAMA mass region. ",
        "conclusions": ""
    },
    "0002/astro-ph0002223_arXiv.txt": {
        "abstract": "We present near--infrared broad--band polarization images of the nuclear regions of the Circinus galaxy in the J, H and K bands.  For the first time the south--eastern reflection cone is detected in polarized light, which is obscured at optical wavelengths behind the galactic disk. This biconical structure is clearly observed in J and H band polarized flux whilst in the K band a more compact structure is detected. Total flux J--K and H--K colour maps reveal a complex colour gradient toward the south--east direction (where the Circinus galactic disk is nearer to us). We find enhanced extinction in an arc shaped structure, at about 200pc from the nucleus, probably part of the star-formation ring.  We model the polarized flux images with the scattering and torus model of Young \\etal, with the same basic input parameters as used by Alexander \\etal in the spectropolarimetry modelling of Circinus. The best fit to the polarized flux is achieved with a torus radius of $\\sim$16pc, and a visual extinction A$_V$, through the torus, to the near--infrared emission regions of $>$66 mags. ",
        "introduction": "\\begin{figure*} \\centerline{\\epsfig{file=figure1a.ps,width=14cm,height=4cm,angle=0}} \\vspace{0.4cm} \\centerline{\\epsfig{file=figure1b.ps,width=14cm,height=4cm,angle=0}} \\vspace{0.4cm} \\centerline{\\epsfig{file=figure1c.ps,width=14cm,height=4cm,angle=0}} \\caption{Near-IR images of the Circinus galaxy, J band (top), H band (middle) and K band (bottom). Contours are arbitrarily scaled.} \\label{} \\end{figure*}  The investigation of nearby Seyfert galaxies with polarimetric observations provides unique information on the nuclear structure of these objects. This information is of particular importance when studying unification models of Seyfert galaxies.  \\begin{figure*} \\centerline{\\epsfig{file=figure2a.ps,width=14cm,height=4cm,angle=0}} \\vspace{0.4cm} \\centerline{\\epsfig{file=figure2b.ps,width=14cm,height=4cm,angle=0}} \\vspace{0.4cm} \\caption{Colour maps of the nuclear regions of the Circinus galaxy. J--K (top) and H--K (bottom). The dark areas correspond to regions of enhanced extinction.} \\label{} \\end{figure*}  The standard unified model for Seyfert galaxies proposes that all types of Seyfert galaxy are fundamentally the same, however, the presence of a dusty molecular ``torus'' obscures the broad line emission in many systems. In this picture the classification of Seyfert 1 or 2 depends on the inclination angle of the torus to the line of sight (Antonucci, 1993). The most convincing evidence for this unified model comes from optical spectropolarimetry. Using this technique, the scattered radiation from the broad line region (BLR) of many Seyfert~2 galaxies is revealed in the form of broad lines in the polarized flux (e.g.\\ Antonucci and Miller, 1985, Young \\etal, 1996a, Heisler, Lumsden and Bailey, 1997).  Near--IR imaging polarimetry provides valuable information on the nature of the polarizing source (e.g.\\ Lumsden \\etal, 1999, Tadhunter \\etal, 1999, Young \\etal, 1996b, Packham \\etal, 1996, 1997, 1998, 1999). For example, in NGC1068, bipolar scattering cones are clearly detected in polarized flux (Young \\etal 1996b, Packham \\etal 1997) and the torus itself has been viewed in silhouette in the H band (Young \\etal 1996b). Interestingly, the structure of the scattering cones often coincide with ground-based narrow band imaging (e.g. Wilson and Tsvetanov, 1994) and high resolution HST imaging (e.g. Capetti \\etal 1997, Falcke, Wilson and Simpson, 1998). Since the unified model infers the presence of a dusty torus obscuring the Seyfert 1 core, we expect that longer wavelength polarimetry will be able to probe deeper into regions within the plane of the torus and to see scattering from the nucleus which might otherwise be shielded from view at optical wavelengths.  Circinus is a nearby (4Mpc) highly inclined (65$\\deg$, Freeman \\etal 1977) Sb-Sd galaxy. At this distance, the spatial scale is 20pc/arcsec. It is seen through a low interstellar extinction window near the Galactic plane (A$_V$ = 1.5 mag; Freeman \\etal 1977). The nuclear optical line ratios are typical of a Seyfert 2 galaxy. This classification as a type 2 is also supported by the detection of intense coronal lines (Oliva \\etal 1994), the intense X-ray Fe 6.4 keV line (Matt \\etal 1996), rapid variation of powerful H$_2$O maser emission and a prominent ionization cone in [O\\,{\\sc iii}]5007 with filamentary supersonic outflows (Marconi \\etal 1994). Recent optical spectropolarimetry (Oliva \\etal, 1998 and Alexander \\etal, 1999a) has shown a scattered polarized broad H$\\alpha$ line and therefore a hidden Seyfert 1 nucleus. Other characteristics include enhanced star forming activity in the form of a star-forming ring of 200pc in size (Marconi \\etal 1994), a Compton thick nucleus at X-ray energies (Matt \\etal 1996) and an unresolved nuclear source ($<$1.5pc) at 2$\\mu$m (Maiolino \\etal 1998).  Since it is already known from optical spectropolarimetry that the Circinus galaxy harbours a hidden type 1 nucleus and the optical ionization cone suggests the presence of an obscuring/collimating torus, this galaxy is an ideal candidate to test the theoretical models of unified schemes. The dusty nature of the Circinus galaxy favours near-IR polarimetry over optical polarimetry due to the lower optical depth at near-IR wavelengths.  In this paper we present near--IR imaging polarimetry revealing, for the first time, a biconical emission region. We investigate the nature of this emission in the context of the standard unified model of Seyfert galaxies.    \\begin{figure*} \\centerline{\\epsfig{file=figure3a.ps,width=14cm,height=5cm,angle=0}} \\vspace{0.4cm} \\centerline{\\epsfig{file=figure3b.ps,width=14cm,height=5cm,angle=0}} \\vspace{0.4cm} \\centerline{\\epsfig{file=figure3c.ps,width=14cm,height=5cm,angle=0}} \\caption{Near-IR polarization maps of Circinus. The J band data is shown at the top, H band in the middle and K band at the bottom. The contours are of the total flux and are arbitrarily scaled. A 1 arcsec polarization vector corresponds to a 10 percent polarization.}  \\label{} \\end{figure*}   \\begin{table} \\begin{minipage}[t]{5.5in} \\caption{Nuclear polarization} \\label{tab:table} \\leavevmode \\footnotesize \\begin{tabular}[h]{cccc} Waveband & aperture & $\\%$ &angle\\\\ &  (arcsec)& & ($\\deg$)\\\\ &\t&                    &             \\\\ &    1   &  2.04$\\pm$0.38    &  34.1$\\pm$3.8\\\\ J        &    2   &  1.71$\\pm$0.16   &  38.3 $\\pm$2.0    \\\\ &    4   &  1.40$\\pm$0.08   &  46.3$\\pm$1.1 \\\\ &\t&                    &             \\\\ &    1      &  2.01$\\pm$0.15 &  33.4$\\pm$2.4   \\\\ H        &    2      &  1.75$\\pm$0.13 &  35.5$\\pm$1.9   \\\\ &     4      &   1.34$\\pm$0.07    & 41.1$\\pm$1.1 \\\\ &\t&                    &             \\\\ &    1       &   3.25$\\pm$0.27  &   32.1$\\pm$3.1 \\\\ K       &    2      &  2.78$\\pm$0.25 &  33.7 $\\pm$2.9  \\\\ &     4    &   1.99$\\pm$0.15 &  34.7 $\\pm$1.8       \\\\ \\end{tabular} \\end{minipage} \\end{table} ",
        "conclusions": "We have presented near--infrared polarimetric images of the Circinus galaxy showing a clear bipolar scattering cone in the J band. The south--east cone was previously undetected at optical wavelengths because it was hidden behind the heavy extinction of the galactic disk. At longer wavelengths, the H and K band images show more compact structures due to the dominance of dichroic absorption over scattered radiation. We have successfully applied an adapted version of the Y95 model to interpret the observed polarized flux distribution. This model includes two geometrical identical scattering cones, diametrically opposite to each other, with the forward cone in the south--east direction at a position angle of 125$\\deg$ and an opening half angle of 45$\\deg$. The inclination of the system to the line of sight is 50$\\deg$. The estimated optical extinction A$_V$ to the nucleus through the torus is $>$66 mag. We estimate that the putative torus in the Circinus galaxy has an outer radius of $\\sim$16pc."
    },
    "0002/astro-ph0002015_arXiv.txt": {
        "abstract": "The variable warm absorber seen with {\\em ASCA} in the X-ray spectrum of MCG$-$6-30-15 shows complex time behaviour in which the optical depth of OVIII anticorrelates with the flux whereas that of OVII is unchanging. The explanation in terms of a two zone absorber has since been challenged by {\\em BeppoSAX} observations. These present a more complicated behaviour for the OVIII edge. We demonstrate here that the presence of a third, intermediate, zone can explain all the observations. In practice, warm absorbers are likely to be extended, multi-zone regions of which only part causes directly observable absorption edges at any given time. ",
        "introduction": "X-ray absorption by partially ionized, optically thin material, along the line of sight of the central engine, the so called {\\em warm absorber}, is one of the prominent features in the X-ray spectrum of many AGN \\cite{reyno97}. The presence of such gas was first postulated in order to explain the unusual form of the X-ray spectrum of QSO MR 2251-178 \\cite{hal84}.  A direct confirmation of the existence of circumnuclear ionized matter came from {\\em ASCA}, which for the first time was able to resolve the OVII and OVIII absorption edges (at rest energies of 0.74 and 0.87 \\keV, respectively) in the X-ray spectra of the Seyfert 1 galaxy MCG$-$6-30-15 \\cite{fab94}. Systematic studies of warm absorbers in Seyfert 1 galaxies with {\\em ASCA} have shown their ubiquity, being detected in half of the sources \\cite{reyno97}.  Warm absorbers were usually described by single zone, photoionization equilibrium models. Under this assumption, simple variability patterns were expected: increasing ionization of the matter when the source brightens. This was not found by Otani et al. \\shortcite{ota96} during their long-look {\\em ASCA} observation of the nearby (z=0.008) Seyfert 1 galaxy MCG$-$6-30-15. The source varied with large amplitude on short time scales ($10^4$ s or so). The depth of the OVII edge, on the contrary, stayed almost constant, while that of the OVIII edge was anticorrelated with the flux. To explain the behaviour of the OVII and OVIII edges, the authors adopted a multizone model in which the OVII and OVIII edges originate from spatially distinct regions. OVII ions in the region responsible for the OVII edge, the outer absorber, have a long recombination timescale (i.e. weeks or more), whereas the OVIII edge arises from more highly ionized material, the inner absorber, in which most oxygen is fully stripped.  The recombination timescale for the OIX ions in the inner absorber is of the order of $10^4 s$ or less. A decrease in the primary ionizing flux is then accompanied by the recombination of OIX to OVIII, giving the observed variation in the OVIII edge depth.  Orr et al. \\shortcite{orr97} raised the possibility of a more complex warm absorber.  During their MCG$-$6-30-15 {\\em BeppoSAX} observation the depth of the OVII edge was marginally consistent with being constant, whereas the optical depth for OVIII, $\\tau(OVIII)$, exhibited significant variability. The authors claimed that its large value during epoch 1\\footnote{The epochs in the {\\em BeppoSAX} observation are chronologically numerated (i.e. number 1 corresponds to the first epoch of the observation, etc.).} ($1.7\\pm0.5$, $1\\sigma$ uncertainty) was inconsistent with the values at all other epochs.  They also pointed out that the value of $\\tau(OVIII)$ during epoch 5 (a low luminosity state) did not show the expected anticorrelaction with the ionizing flux .  In both Otani et al. \\shortcite{ota96} and Orr et al.\\shortcite{orr97}, $\\tau(OVIII)$ was plotted versus count rate.  In order to compare {\\em ASCA} and {\\em BeppoSAX} observations, these two observational results for $\\tau(OVIII)$ have been plotted versus luminosity\\footnote{The conversion from count rate to luminosity has been obtained using $H_o=50$ \\kmpspMpc, $q_0=0$ and PIMMS (http://heasarc.gsfc.nasa.gov/Tools/w3pimms.html).} in 0.1-10 \\keV in figure 1 (instead of versus count rate). The conversion from count rate to luminosity is different for each apparatus and therefore, the comparison can not be made using count rate, but luminosity. \\begin{figure} \\centerline{\\psfig{figure=fig1.ps,width=0.50\\textwidth,angle=-90}} \\caption{Comparison {\\em ASCA} (square) - {\\em BeppoSAX} (circle) observational results for the optical depth of OVIII, $\\tau(OVIII)$, versus luminosity in 0.1-10 \\keV. Note the only discrepant point in the {\\em BeppoSAX} observation labelled as 1 and the very good agreement for point 5.} \\end{figure} The only discrepancy between both sets of data is the result for epoch 1 (labelled 1 in figure 1)\\footnote{This epoch was at the very beginning of the {\\em BeppoSAX} observation. In a private communication, A. Orr notes that this high value for the optical depth for OVIII could not be due to instrumental effects, since no similar behaviour, neither at the beginning nor during the observation has been observed in any other {\\em BeppoSAX} target yet.}. Note also that there is no disagreement at all with any other point (Orr et al. mentioned epoch 5 as problematic, but in this plot it appears consistent with the rest of the Otani et al. data).  In this paper we present a simple photoionization model that accounts for the experimental results of both {\\em ASCA} and {\\em BeppoSAX} observations. Section 2 describes the code used to model the inner warm absorber. The application to the MCG$-$6-30-15 warm absorber is presented in Section 3. The extension of the multi-zone model for the warm absorber (i. e. the third component) is addressed in Section 4. Finally our conclusions are discussed in Section 5. ",
        "conclusions": "Warm absorbers have been the subject of extensive studies during the last decade. Such regions are not spatially resolved, and all the available information about their geometry is obtained from analysis of the variability of the oxygen edges. The explanation we offer for the time variability of the MCG$-$6-30-15 warm absorber during both {\\em ASCA} and {\\em BeppoSAX} observations does not invoke complex processes, but a very simple photoionization model together with the presence of a multi-zone warm absorber. This would be constituted by a continuum of clouds at different radii and different densities, such that only some of them contribute to the total optical depth for OVIII depending on the value of the luminosity.  As a final remark note how in our model WA3 is much more volume filling than WA1 ($\\Delta R/R\\geq 0.1$ for WA3 while $\\Delta R/R\\leq 10^{-3}$ for WA1). This suggests that WA3 could be part of the inter-cloud medium of WA1."
    },
    "0002/astro-ph0002153_arXiv.txt": {
        "abstract": "We re-analyze the question of the use of cosmological observations to infer the present state and future evolution of our patch of the universe. In particular, we discuss under which conditions one might be able to infer that our patch will enter an inflationary stage, as a {\\em prima facie} interpretation of the Type Ia supernovae and CMB data would suggest. We then establish a `physical' criterion for the existence of inflation, to be contrasted with the more `mathematical' one recently proposed by Starkman {\\em et al.} \\cite{STV}. ",
        "introduction": "\\label{intro} The issue of the present state, future dynamics and final fate of the Universe, or at least our patch of it, has been recently pushed to the front line of research in cosmology. This is mostly due to observations of high redshift type Ia supernovae, performed by two independent groups (the ``Supernova Cosmology Project'' and the ``High-Z Supernova Team''), which allowed accurate measurements of the luminosity-redshift relation out to redshifts up to about $z \\sim 1$ \\cite{Perlmu,Riess1,Garnavich}. It should be kept in mind that these measurements are done on the assumption that these supernovae are standard candles, which is by no means demonstrated and could conceivably be wrong. There are concerns about the evolution of these objects and the possible dimming caused by intergalactic dust \\cite{Drell,Riess2}, but we will ignore these for the purposes of this paper, and assume that the quoted results are correct.  The supernovae data, when combined with the ever growing set of CMBR anisotropy observations, strongly suggest an accelerated expansion of the Universe at the present epoch, with cosmological parameters $\\Omega_\\Lambda \\sim 0.7$ and $\\Omega_{\\rm m} \\sim 0.3$. A further cause of concern here is the model dependence of the CMBR analysis, but we shall again accept the above results for the purpose of this paper.  Taken at face value, these results would seem to show that the universe will necessarily enter an inflationary stage in the near future. However, as pointed out by Starkman, Trodden and Vachaspati \\cite{STV} this is not necessarily so. We could be living in a small, sub-horizon bubble, for example. And even if we were indeed inflating, it would not be trivial to demonstrate it. In the above work, these authors looked at the crucial question of `How far out must we look to infer that the patch of the universe in  which we are living is inflating?' Their analysis is based on previous work by Vachaspati and Trodden \\cite{VT} which shows that the onset of inflation can in some sense be identified with the comoving contraction of our minimal anti-trapped surface (MAS)\\footnote{The MAS of each comoving observer is a sphere centered on him/her, on which the velocity of comoving objects is $c$. For the particular case of an homogeneous universe, the MAS has a physical radius $cH^{-1}$.}. They then argue that if one can confirm cosmic acceleration up to a redshift $z_{MAS}$ and detect the contraction of our MAS, then our universe must be inflating. Unfortunately, even if we can do the former (for $\\Omega_\\Lambda=0.8$ the required redshift is $z_{MAS}\\sim1.8$), it turns out that there is no way to presently confirm the latter, because the accelerated expansion hasn't been going for long enough for the MAS to contract. Only if we had $\\Omega_\\Lambda\\ge0.96$ would we be able to demonstrate inflation today.  As in the proverbial mathematicians joke, the method outlined by Starkman {\\em et al.} provides an answer that is completely accurate but will take a long time to find, and hence is of no immediate use to us. In this paper, however, we will explore a different possibility. Our main aim is to provide what could be called a physicists version of the ``mathematical'' question of Starkman {\\em et al.} \\cite{STV}. In other words, we are asking, {\\em `If we can't know for sure the fate of the universe at present, what is our best guess today?'}. As we will discuss, we can answer this question, although it will involve making some crucial additional assumptions.  In order to answer the above question we compare the particle and event horizons. We show that for a flat universe with $\\Omega_\\Lambda \\gsim 0.14$ the particle horizon is greater than the distance to the event horizon meaning that today we may be able to observe a larger portion of the Universe than that which will ever be able to influence us. We argue that if we find evidence for a constant vacuum density up to a distance from us equal to the event horizon then our Universe will necessarily enter an inflationary phase in the not too distante future, {\\em assuming} that the potential of the scalar field which drives inflation is time-independent and that the content of the observable universe will remain `frozen' in comoving coordinates.  Note that Starkman {\\em et al.} argue that inflation can only take place if the vacuum energy dominates the energy density on a region with physical radius not smaller than that of the MAS at that time. However, they did not assume that the content of the observable universe would remain frozen in comoving coordinates and so they found that the larger is the contribution of a cosmological constant for the total density of the Universe, the larger is the redshift out to which one has to look in order to infer that our portion of the universe is inflating. This result seems paradoxical, until one realizes that the size of the MAS at a given time {\\em does not}, by itself, say anything about inflation. The main reason why $z_{MAS}$ grows with $\\Omega_\\Lambda$ is simply because the scale factor has grown more.  The plan of this paper is as follows. In the next section we introduce the various lengthscales that are relevant to our discussion, and provide a qualitative discussion of our test for inflation. We also discuss the assumptions involved and compare our `physical' test with the `mathematical' one recently proposed by Starkman {\\em et al.} \\cite{STV}. In section \\ref{biaf} we provide a more quantitave analysis of our criterion. We also discuss in more detail our crucial assumption of an energy-momentum distribution which remains frozen in comoving coordinates. Finally, in section \\ref{conc} we summarize our results and discuss some other outstanding issues. ",
        "conclusions": "\\label{conc} We have provided a simple analysis of the use of cosmological observations to infer the state and fate of our patch of the universe. In particular, in the same spirit of Starkman {\\em et al. } \\cite{STV}, we have discussed possible criteria for inferring the present or future existence of an inflationary epoch in our patch of the universe.  We have presented a `physical' criterion for the existence of inflation, and contrasted it with the `mathematical' one that has been introduced in \\cite{STV}. Ours has the advantage of being able to provide (in principle) a definite answer at the present epoch, but the disadvantage of ultimately relying on assumptions on the content of the local universe and on field dynamics. We consider our assumptions to be plausible, but we can certainly conceive of (arguably contrived or fine-tuned) mechanisms that would be capable of violating it."
    },
    "0002/astro-ph0002479_arXiv.txt": {
        "abstract": "For the low density filamentary and sheet-like structures in the warm ($\\sim 10^4$ to $\\sim 10^6$ K) IGM predicted by numerical simulations the resonant line scattering of X--ray background (XRB) photons by  He and H--like ions of heavy elements can exceed the ``local'' thermal emission by a factor of a few or more. Due to the conservative nature of scattering this resonantly scattered radiation can only be identified  if a significant fraction of the XRB is resolved and removed. While the combined spectrum of the resolved sources will contain X--ray absorption features, the residual background will contain corresponding emission features with the same intensity. At the relevant  densities and temperatures  the lines of  He and H--like oxygen at 0.57 and 0.65 keV are most promising. These lines (which have a typical width of $\\sim$ 1--2 eV) may contain up to 50\\% of the total 0.5--1 keV emission of the filament. For a nearby ($z \\la 0.1$) filament with a Thomson optical depth of $10^{-4}$  XMM should detect about 200 photons in the OVII line during a $10^5$ s exposure if the metallicity of the gas is as large as observed in galaxy clusters. On average up to a few percent of  the soft XRB  could be resonantly scattered  by this phase of the IGM  and resonantly scattered photons should account for a significant fraction of the truly diffuse background at low energies. Close to bright  X-ray sources like galaxy  clusters or AGN the flux of scattered radiation will be further enhanced. Off-line blazars are the most promising illuminating sources. The scattered emission from AGN may also constrain the duration of the active phase of these objects. ",
        "introduction": "At high redshift the majority of baryons are  contained in a photoionized IGM with temperature $\\sim 10^4\\K $  which is responsible for the strong  \\op absorption in the spectra of high redshift  objects (see Rauch 1998 and Weinberg 1999 for reviews). At low redshift, however, there are little observational constraints on  the thermal state and baryon content of the IGM (e.g. Barcons, Fabian \\& Rees 1991). Numerical simulations predict that a considerable fraction of all baryons is still contained in a warm IGM which traces the filamentary and sheet-like distribution of the dark matter (e.g. Ostriker \\& Cen, 1996, Cen \\& Ostriker 1999). This phase of the IGM should have densities of a few up to a few tens times the  mean baryonic density. The temperatures should range between a `` photoionization temperature'' at low densities ($\\sim 10^3 -10^4\\K$) and the virial temperature of sheets and filaments (~$\\sim 10^6\\K$). These temperatures might be enhanced due to the energy  input by star formation and AGN and it has been argued that such an energy input is  necessary to explain  the X-ray luminosity temperature relation  of galaxy clusters (Kaiser 1991; Ponman, Cannon \\& Navarro 1999; Pen 1999). This ``warm'' phase of the IGM is of special interest as it will contain a record of the energy and metals expelled from galaxies. The typical surface brightness  of filaments due to local thermal emission of the diffuse gas is, however, a factor 100 or more smaller than that  of the X-ray background. This local thermal emission will be detectable  with XMM only for especially strong filaments (Pierre, Bryan \\& Gastaud 1999) but see Scharf et al. (1999) for a claimed detection of such a filament with ROSAT. It was also suggested that the warm IGM  produces measurable absorption in the resonant transitions of heavy elements such as oxygen or iron if there is a bright quasar  behind the gas (Shapiro and Bahcall 1980, Aldcroft et al., 1994, Hellsten, Gnedin, Miralda--Escude 1998, Perna and Loeb, 1998, Markevitch, 1999). This seems a promising method to study  the warm IGM, especially as long as  high energy resolution is only possible for bright sources (i.e. using gratings). It will, however, only give information along the line-of-sight  to point  sources. We explore here  the  possibility to investigate the warm IGM by its emission due to resonant scattering of X-ray background photons  by He and H--like ions of heavy elements. While for the  high density and  temperatures in galaxy clusters the local thermal emission   clearly dominates over resonant scattering the opposite is true for the more moderate densities and temperatures expected in the sheet-like and filamentary structures of the warm phase of the IGM.  In section 2 we discuss the importance of resonant scattering relative to the local thermal emission due to collisional excitation and ionization for the temperatures and densities prevalent in the warm IGM. Section 3 discusses the detectability of filamentary structures with upcoming satellite missions. In section 4 we discuss the illumination of filamentary structures by individual bright sources and section 5 contains our conclusions.  A Hubble constant of $H_0=50~\\kms~\\Mpc^{-1}$ and an Einstein-de-Sitter Universe was assumed throughout the paper. For the baryon density we take $\\rho_{\\rm bar}= 3.6~10^{-31}~\\g~\\cm^{-3}$ ($\\Omega_{\\rm bar} h^2 = 0.02$,  Burles et al., 1999). The abundances of heavy elements were assumed to be a constant fraction of  solar abundances as given by Feldman (1992). For the local  XRB spectrum we use the simple approximation $I(E)=I_0\\, E^{-1.3} e^{-\\frac{E}{40\\keV}}$ for  energies $E \\ge 1 \\keV$ and a power law with a photon index of 2 below 1 keV, where  $I_0 = 8~\\phot~\\cm^{-2}~\\sec^{-1}~\\keV^{-1}~\\sr^{-1}$ (e.g. Barcons \\& Fabian 1992, Miyaji et al. 1998). ",
        "conclusions": "We have demonstrated that for the temperatures and densities prevalent in the the filamentary and sheet-like structures of the warm intergalactic medium illumination by  XRB photons significantly increases the X--ray emissivity of the medium. The strongest contribution to the X--ray emissivity is due to resonant scattering of XRB photons. This resonantly scattered emission can be detected in images from  which  a significant fraction of the XRB due to compact sources has  been removed.  Resonant lines of He and H-like oxygen are most promising for detecting filamentary structures in the warm IGM. The fraction of oxygen in the ionization states \\ovii\\ or \\oviii\\ is  larger than 30 percent practically for the whole range of densities  and temperatures expected in the warm IGM.  Estimates for the detectability of filaments in the warm IGM  which take only local thermal emission into account may have been overly pessimistic. XMM should be able to detect the radiation resonantly  scattered by diffuse gas in a filament with Thomson optical depth of $10^{-4}$ and  metallicity 0.3 solar at $z\\sim 0.1$ with signal-to-noise of a few.  For  filaments close to massive cluster the signal will  be enhanced  due to illumination from the cluster core and/or AGN contained in the cluster. The most promising sources emitting high X-ray intensities are beamed blazars which should have a high space density.  For AGN and especially blazars constraints on the duration of the active phase may also be obtainable."
    },
    "0002/astro-ph0002129_arXiv.txt": {
        "abstract": "We present the data and analysis of VLBI observations at 1.6, 5 and 15 GHz of a sample of faint Gigahertz Peaked Spectrum (GPS) sources selected from the Westerbork Northern Sky Survey (WENSS).  The 5 GHz observations involved a global array of 16 stations and yielded data on the total sample of 47 sources.  A subsample of 26 GPS sources with peak frequencies $\\nu_p > 5$ GHz and/or peak flux densities $S_p > 125 $ mJy was observed with the VLBA at 15 GHz. A second subsample of 29 sources, with $\\nu_p <5$ GHz, was observed at 1.6 GHz using a 14 station global VLBI array. In this way, 44 of the 47 sources (94\\%) in the sample were observed above and at or below their spectral peak. Spectral decomposition allowed us to identify 3, 11, 7, and 2 objects as compact symmetric objects, compact doubles, core-jet and complex sources respectively. However, many of the sources classified as compact double or core-jet sources show only two components making their classification rather tentative. This may explain why the strong morphological dichotomy of GPS quasars and galaxies found for radio-bright GPS sources, is not as clear in this faint sample. ",
        "introduction": "Gigahertz Peaked Spectrum (GPS, e.g. O'Dea 1998) are a class of extragalactic radio source, characterised by a convex shaped radio spectrum peaking at about 1 GHz in frequency, and sub-galactic sizes. Their small sizes make observations using Very Long Baseline Interferometry (VLBI) necessary to reveal their radio morphologies. Early VLBI observations showed that some GPS sources identified with galaxies have Compact Double (CD) morphologies (Philips and Mutel, 1982), and it was suggested that these were the mini-lobes of very young or alternatively old, frustrated objects (Philips and Mutel, 1982; Wilkinson et al. 1984, van Breugel, Miley and Heckman, 1984). Later, when reliable VLBI observations at higher frequencies became possible, it was found that some of the CD-sources had a compact flat spectrum component in their centres (Conway et al. 1992, Wilkinson et al. 1994). These flat spectrum components were interpreted as the central cores, and many CD-sources were renamed compact triples or Compact Symmetric Objects (CSO, Conway et al. 1992, Wilkinson et al. 1994). High dynamic range VLBI observations by Dallacasa et al (1995) and Stanghellini et al. (1997) have shown that most GPS galaxies indeed have jets leading from the central compact core to the outer hotspots or lobes. This is in contrast to the GPS sources identified with quasars, which tend to have core-jet morphologies with no outer lobes (Stanghellini et al. 1997). Snellen et al. (1999) have shown that the redshift distributions of the GPS galaxies and quasars are very different, and that it is therefore unlikely that they form a single class of object unified by orientation. They suggest that they are separate classes of object, which just happen to have the same radio-spectral morphologies.  The separation velocities of the hotspots have now been measured for a small number of GPS galaxies to be $0.2h^{-1}$c (Owsianik and Conway, 1998; Owsianik, Conway and Polatidis, 1998; Tschager et al. 1999). This makes it very likely that these are young objects of ages typically $\\sim 10^3$ yr (assuming a constant separation velocity), rather than old objects constrained in their growth by a dense ISM. These are therefore the objects of choice to study the early evolution of extragalactic radio sources.  In the past, work has been concentrated on samples of the radio brightest GPS sources (eg. O'Dea et al 1991). In order to disentangle radio power and redshift effects on the properties of GPS sources, we constructed a sample of faint GPS sources from the Westerbork Northern Sky Survey (WENSS, Rengelink et al. 1997), which in combination with other samples allows, for the first time, the study of these objects over a large range of flux density and radio spectral peak frequency. The construction of the faint sample is described in Snellen et al. (1998a); the optical and near-infrared imaging is described in Snellen et al. (1998b); and the optical spectroscopy in Snellen et al. (1999a). This paper describes multi-frequency VLBI observations of the sample, and the  radio-morphologies of the individual sources. What can be learned from the faint GPS sample about radio source evolution is discussed in the accompanying paper (Snellen et al. 2000). ",
        "conclusions": "Multi-frequency VLBI observations have been presented of a faint sample of GPS sources. All 47 sources in the sample were successfully observed at 5 GHz, 26 sources were observed at 15 GHz, and 20 sources were observed at 1.6 GHz. In this way 94\\% of the sources have been mapped above and below their spectral peak. The spectral decomposition allowed us to classify 3 GPS galaxies as compact symmetric objects (CSO), 1 galaxy and 1 quasar as complex (CX) sources, 2 quasars and 5 galaxies as core-jet (CJ) sources, and 9 galaxies and 2 quasars as compact doubles (CD). Twenty-five of the sources could not be classified, 20 because they were too compact. The strong morphological dichotomy of GPS galaxies and quasars found by Stanghellini et al. (1997) in their radio bright GPS sample is not so clear in this sample. However, many of the sources classified as CD and CJ have a two-component structure, making their classification only tentative."
    },
    "0002/astro-ph0002403_arXiv.txt": {
        "abstract": "In a series of papers we investigate far-infrared emission from rich galaxy clusters. Maps have been obtained by ISO at $60\\mu {\\rm m}$, $100\\mu {\\rm m}$, $135\\mu {\\rm m}$, and $200\\mu {\\rm m} $ using the PHT-C camera. Ground based imaging and spectroscopy were also acquired. Here we present the results for the cooling flow cluster S\\'{e}rsic\\,159-03. An infrared source coincident with the dominant cD galaxy is found. Some off-center sources are also present, but without any obvious counterparts. ",
        "introduction": "\\label{introduction} The first paper in this series (Hansen et al. \\cite{han99}, paper{\\sc \\,i}) presented infrared data for the Abell\\,2670 cluster. We identified 3 far-infrared sources apparently related to star forming galaxies in the cluster. The present paper concerns the rich cluster S\\'{e}rsic\\,159-03. The central part of the S\\'{e}rsic\\,159-03 cluster was mapped by the Infrared Space Observatory (ISO) satellite, using the PHT-C camera (Lemke et al. \\cite{lem96}) at $60\\mu {\\rm m}$, $100\\mu {\\rm m}$, $135\\mu {\\rm m}$, and $200\\mu {\\rm m} $. The observations were performed twice with slightly different position angles which gives an opportunity to do independent detections and to study possible instrumental effects.  The S\\'{e}rsic\\,159-03 cluster (Abell S1101, z=0.0564) is of richness class 0, Bautz-Morgan type{\\sc \\,iii} with a central dominant cD galaxy (Abell et al. \\cite{abe89}). A cooling flow is present, and Allen and Fabian (\\cite{all97}) found a mass deposition rate of ${\\rm \\dot{M} = 231^{+11}_{-10} ~M_{\\sun} ~yr^{-1}}$ from ROSAT PSPC data. The cooling flow is centered on the cD~galaxy which exhibits nebular line emission. Crawford and Fabian (\\cite{cra92}) obtained optical spectra and found from line-ratio diagrams that the ratios obtained along the slit bridged the gap between class{\\sc \\,i} and class{\\sc \\,ii} in the scheme of Heckman et al.  (\\cite{hec89}). Their spectra had position angle $90\\degr$. West of the center they discovered a detached filament of emission having extreme class{\\sc \\,ii} characteristics. They argued that the different line ratios are due to changes in ionization properties. Below in Fig.~\\ref{emission} we show the  extent of the nebular emission. In a subsequent paper  Crawford and Fabian (\\cite{cra93}) included IUE data to obtain the optical-ultraviolet continuum. They announced that a strong Ly$\\alpha$ line is present in the IUE spectrum. ",
        "conclusions": "\\label{conclusion} The availability of two observations covering essentially the same field at several wavelengths allows us to identify 4 faint ($\\approx 0.1$~Jy) far-infrared sources with some confidence. A central source, C100-1, is attributed to the cD galaxy which contains optical filaments, but our optical images do not reveal significant evidence of dust lanes. The fluxes measured for C100-1 are of the same order of magnitude as expected from dust related to star formation in the cooling flow. For the non-central sources we cannot point out any particular optical candidates in contrast to the results from the Abell\\,2670 field (paper{\\sc \\,i}) where galaxies with enhanced star formation were found coincident with the infrared sources."
    },
    "0002/astro-ph0002259_arXiv.txt": {
        "abstract": "The cold diffuse interstellar medium must harbor pockets of hot gas to produce the large observed abundances of molecular species, the formation of which require much more energy than available in the bulk of its volume. These hot spots have so far escaped direct detection but observations and modeling severely constrain their phase-space structure \\ie{} they must have a small volume filling factor (a few \\%), surface filling factors larger than unity with large fluctuations about average and comparable velocity structure in ``pencil beams'' and ``large beams''.  The dissipation of the non-thermal energy of supersonic turbulence occurs in bursts, either in shocks or in the regions of large shear at the boundary of coherent vortices. These two processes are susceptible to generate localized hot regions in the cold medium.  Yet, it is of interest to determine which of them, if any, dominates the dissipation of turbulence in the cold interstellar medium.  In this paper, we analyze the spatial and kinematic properties of two subsets of hydrodynamical compressible turbulence: the regions of largest negative divergence and those of largest vorticity and confront them with the observational constraints. We find that these two subsets fulfill the constraints equally well. A similar analysis should be conducted in the future on simulations of MHD turbulence. ",
        "introduction": "For several decades now, molecules have been detected in the diffuse component of the cold neutral medium and these observations raise several intriguing questions.  Firstly, the large observed abundances of \\CHp{} \\cite{crane95:hrsiCH,gredel97:iCHsOBa} and \\HCOp{} and OH \\cite{lucas96:pbsgHCOatecs} in the cold diffuse medium imply that activation barriers and endothermicities of several thousands Kelvin be overcome.  The formation of \\CHp{} proceeds through the endothermic reaction of \\Cp{} with \\HH{} ($\\Delta E/k$=4640 K) and that of OH via the reaction of O with \\HH{} which has an activation energy $\\Delta E/k$ =2980 K. In the diffuse gas, \\HCOp{} forms from \\CHthp, a daughter molecule of \\CHp. None of the large observed abundances can be explained by standard steady-state chemistry in cold diffuse gas.  Pockets of hot gas must therefore exist in the cold diffuse medium.  Secondly, many other molecules like CO, CS, SO, CN, HCN, HNC, H$_2$S, C$_2$H have now been detected in absorption in front of extragalactic continuum sources in local or redshifted gas (Lucas \\& Liszt 1993, 1994, 1997; Liszt \\& Lucas 1994, 1995, 1996; Wiklind \\& Combes 1997, 1998). \\nocite{lucas93:pbommwmaBLL,lucas94:pbommwma,lucas97:mwodc} \\nocite{liszt94:mHCOedctdo,liszt95:mwmeaca,liszt95:mwmeaca} \\nocite{wiklind97:malhr,wiklind98:cmals} The lines of sight sample the edges of molecular clouds or the diffuse medium, which corresponds to gas poorly shielded from the ambient UV radiation field.  This derives from the low excitation temperatures (often close to the temperature of the cosmic background) measured in these transitions.  Yet, the molecular abundances derived are close to those of dark clouds.  Thirdly, the spatial and velocity distribution of these molecules is highly elusive. Observations of CO emission and absorption lines toward extragalactic sources by Liszt \\& Lucas \\cite*{liszt98:COaecercs} show that: {\\it (i)} absorption and emission line profiles have comparable linewidths whereas the projected size of the sampled volumes of gas differ by more than four orders of magnitude (60 vs. 10$^{-3}$ arcsec), {\\it (ii)} the excitation temperature of the CO molecules changes only weakly across the profiles and {\\it (iii)} there are very few lines of sight with no absorption line detected. The phase-space distribution of the CO--rich gas in diffuse clouds must therefore have a surface filling factor close to unity, and its velocity field must be as adequately sampled by a pencil beam than by a large beam. As mentioned by the authors, these observations suggest a one-dimensional structure for the molecular component, which contrasts with the profuse small scale structure observed in emission.  Two mechanisms, operating at very different size scales, can trigger a {\\it hot chemistry} in the cold diffuse medium and produce the observed molecular abundances: magneto-hydrodynamical (MHD) shocks~\\cite{elitzur78,draine86:mhdsdc:ps,draine86:mhdsdc:lstro,pineau86:tsisms:fdc,flower98:ctsim:pal} and intense coherent vortices, responsible for a non-negligible fraction of the viscous dissipation of supersonic turbulence \\cite{falgarone95:iticigktdhpfgm,falgarone95:csitldic,joulain98:necdsit}. This is so because the dissipation of the non-thermal energy of supersonic turbulence is concentrated in shocks and in the regions of large shear at the boundary of coherent vortices. In both cases, only a few \\% of hot gas on any line of sight across the cold medium is sufficient to reproduce the observed column densities of molecules. This small fraction of hot gas corresponds to about 6 MHD shocks of 10 \\kms{} or 1000 vortices with rotation velocity of 3.5~\\kms{} per magnitude of gas (or $\\Ntot=1.8\\times 10^{21}$ \\cq) of density $n \\approx 30$ \\cc{} \\cite{verstraete99:hcdm:ssHHrl}.  The main problem met with the models of individual C shocks or vortices is that the predicted shift between neutral and ionized species is larger than observed.  In this paper, we simply address the issue of the impact of the line of sight averaging upon the resulting line profiles in a turbulent velocity field. We investigate the spatial and velocity distribution of the regions of high vorticity or high negative divergence (shocks) in a simulation of compressible turbulence to test whether or not the space-velocity characteristics of any of these subsets fulfill the observational requirements.  We carry our analysis on the numerical simulations of compressible turbulence of Porter, Pouquet \\& Woodward \\cite*{porter94:ksdtdsf}. They are $512^3$ simulations of midly compressible turbulence (initial rms Mach number=1). The time step analysed here is $t=1.2 \\Tac$ where \\Tac{} is the acoustic crossing time. At that epoch in the simulation, many shocks have survived producing density contrasts as large as 40, but the bulk of the energy at small scales is contained in the vortical modes.  These simulations are hydrodynamical and do not include magnetic field. The impact of magnetic field upon the statistical properties of turbulence may be less important than foreseen.  Recent simulations of MHD compressible turbulence \\cite{ostriker99:ksesgmc:ddt} show that the dissipation timescale of MHD turbulence is closer than previously thought to that of hydrodynamical turbulence.  Descriptions of the energy cascade in models of MHD turbulence also predict that vortices must play an important role~\\cite{goldreich95:tist.II:sat,lazarian99:rwsf}. Magnetic field does not prevent either the intermittency of the velocity field to develop \\cite{brandenburg96,galtier98:imhdf,politano98:khemhdctolscf,politano98:dlstmf,politano95:mimhdt}. In MHD as in hydrodynamic turbulence, shocks interact and generate vortex layers, which are Kelvin-Helmholtz unstable and eventually form vorticity filaments.  Another limitation of our study lies in the fact that the Reynolds number of the simulations is small compared to that of interstellar turbulence. Recent studies of high Reynolds number turbulence in Helium bring the unexpected result that the statistical properties of the velocity field have little dependence with the Reynolds number, at large Reynolds numbers~\\cite{tabeling96:pdfsflrnt}. We therefore believe that the analysis presented here on hydrodynamical simulations has some relevance to the understanding of interstellar turbulence.  We discuss the velocity structure of the two subsets (regions of high vorticity or of high negative divergence) and compare in each case the velocity samplings provided by pencil beams versus large beams (Sect.~\\ref{sec:vssivsct}). In Sect.~\\ref{sec:sdhvds}, we discuss the spatial distributions of the two subsets. In Sect.~\\ref{sec:co}, we compare the results derived from the numerical simulations with the observations. ",
        "conclusions": "The regions of largest vorticity or largest divergence in compressible turbulence are small subsets of a whole turbulent velocity field but we have shown that, despite their small volume filling factor (here $\\fvol=0.03$):  (i) they sample the whole velocity field,  (ii) pencil beam samplings across these subsets have velocity coverages almost as broad as those obtained with large beams. The values (35\\% to 45\\%) and the signs of these differences are consistent with the observations.  (iii) the PDFs of centroid velocity increments built on these subsets have more extended non-Gaussian wings than those of the full velocity field. Their intermittent characteristics subsist at larger lags than for the full velocity field.  (iv) copious small scale structure with large contrasts is seen in the maps of the surface filling factor of the regions of large vorticity and large divergence.  These contrasts are of the same order of magnitude as those observed in the column densities of \\CHp{} for instance for a given column density of sampled gas.  (v) the surface filling factor of the subsets of high vorticity/divergence are almost everywhere larger than unity even though their volume filling factor is as small as 3\\% in the subsets studied here. Vortices are slightly more efficient than shocks at covering the sky: they tend to be more numerous, and to form smaller structures which are less clustered in space.  In summary our study confirms that mild shocks as well as intense vortices could be the subsets of the cold diffuse medium enriched in molecules and responsible for the molecular absorption lines detected in the direction of extragalactic sources.  Numerical simulations of hydrodynamical turbulence are not ideally suited to test the properties of such subsets, but one-fluid simulations of MHD turbulence are not ideal either because of the importance of the ion-neutral streaming in the triggering of hot chemistry, whether in MHD shocks or in magnetized vortices. More detailed predictions of the phase-space structures of the subsets of turbulence where hot chemistry is activated in the cold diffuse medium require calculations of MHD compressible turbulence which take ion-neutral drift into account."
    },
    "0002/astro-ph0002290_arXiv.txt": {
        "abstract": "We have calculated the evolution of low metallicity red giant stars under the assumption of deep mixing between the convective envelope and the hydrogen burning shell.  We find that the extent of the observed abundance anomalies, and in particular the universal O-Na anticorrelation, can be totally explained by mixing which does not lead to significant helium enrichment of the envelope. On the other hand, models with extremely deep mixing and strong helium enrichment predict anomalies of sodium and oxygen, which are much larger than the observed ones. This latter result depends solely on the nucleosynthesis inside the hydrogen burning shell, but not on the details of the mixing descriptions. These, however, influence the evolution of surface abundances with brightness, which we compare with the limited observational material available.  Our models allow, nevertheless, to infer details on the depth and speed of the mixing process in several clusters. Models with strong helium enrichment evolve to high luminosities and show an increased mass loss. However, under peculiar assumptions, red giants reach very high luminosities even without extreme helium mixing. Due to the consequently increased mass loss, such models could be candidates for blue horizontal branch stars, and, at the same time, would be consistent with the observed abundance anomalies. ",
        "introduction": "The observed anomalies in CNO-, NeNa- and MgAl- elements in globular cluster red giants (see \\cite{kra:94} and \\cite{daco:98} for reviews) are unexplained in canonical low-mass star evolution theory and indicate effects beyond the standard picture. At least for anomalies in those isotopes participating in the CNO-cycle models relying on the assumption of an additional, non-standard mixing process inside the stars have been presented, which explain convincingly the observations, including the evolution of the carbon abundance along the RGB, i.e.\\ with time (see \\cite{cchar:95}; \\cite{dw:96}; \\cite{csb:98}). This mixing is supposed to set in after the hydrogen burning shell has reached the composition discontinuity left behind by the first dredge-up on the red giant branch (RGB), i.e., after the so-called RGB bump (e.g. \\cite{sm:79}; \\cite{cchar:95}; \\cite{cbw:98}) where the molecular weight barrier between hydrogen burning shell and envelope is at a minimum.  It is usually described in terms of a diffusion process of certain efficiency and penetration depth.  The CNO anomalies and their correlations with brightness and metallicity (which are expected qualitatively from nucleosynthesis arguments in standard models, as discussed by \\cite{csb:98}) are thus explained by a purely evolutionary picture (\\cite{st:92}; \\cite{dw:96}; \\cite{bs:99b}).  The physical origin of the mixing process is believed to be found in differential rotation of the star and the parameters used in some of the presently available calculations have been derived from existing theories (e.g.\\ \\cite{zahn:92}; \\cite{mz:98}), which are, however, far from being complete.  A similar situation holds for oxygen and sodium, which are found to be anti-correlated in globular cluster red giants (\\cite{ksn:93}). \\citen{dd:90} and \\citen{lhs:93} showed that this could result from the mixing of elements participating in the ONeNa-cycle, which operates at higher temperatures than the CNO-cycle and therefore requires deeper mixing.  \\citen{dw:96} demonstrated how all anomalies of the mentioned elements known at that time can be explained by the deep mixing scenario. Their calculations were done by using canonical red giant models, which evolved along the RGB without any mixing, and performing the mixing and nuclear reactions in a post-processing way. For this approach to be correct it is necessary that the evolution of the background models is not affected by the mixing process. In fact, the mixing necessary to reproduce the observed anomalies was always so shallow that only very small amounts of hydrogen/helium were mixed between envelope and shell, even in the case of the O-Na-anti-correlation. This was taken as sufficient justification of the underlying basic assumption. Very similar conclusions were obtained by \\citen{csb:98}.  Within this approach, one cannot investigate the possibility or necessity for even deeper mixing, which would affect the hydrogen/helium structure of the models.  \\citen{swei:97} has renewed the interest in such deep mixing by connecting the problem of horizontal branch morphology with that of observed anomalies, as previously suggested by \\citen{lh:95}.  If the mixing leads to severe helium enhancement in the envelope, increased luminosities and stellar winds result, such that the star will populate the blue horizontal branch (HB), while it remains a red HB star without the additional mixing.  While \\citen{dw:96} did not investigate the effect of helium transport, \\citen{swei:97} did not follow the evolution of the participating isotopes to compare with observations. In the present paper, we therefore attempt to close this gap by computing full evolutionary sequences which include deep diffusive mixing and by investigating abundance anomalies using these self-consistent models as background models. In particular, we want to answer the question {\\em how much helium enrichment of the envelope is necessary or allowed to achieve or to keep consistency with observations.}  In Sect.~2 we will discuss the nucleosynthesis aspects of the problem and review the observational status of the global O-Na anticorrelation, which is a powerful tracer of the transport processes in the red giants. In Sect.~3 we will present and discuss the evolutionary models. After that, the predictions for the abundances based on our mixed models and post-processing nucleosynthesis will follow, before the conclusions close the paper. ",
        "conclusions": "If the O-Na-anticorrelation observed in many globular cluster red giants is indeed due to a deep mixing process beyond the standard effects taken into account in canonical stellar evolution theory, the question is justified whether this deep mixing might affect the H-He-profile as well. In this case, consequences for the red giant evolution including phases of enhanced luminosities and mass loss could result. We have, therefore, discussed both the evolutionary and nucleosynthetic effects quantitatively by performing stellar evolution calculations including deep mixing and post-processing nucleosynthesis models (the latter as we did in our earlier papers \\cite{dw:96}; \\cite{ddn:98}).  {From} arguments depending only on nucleosynthesis we could already infer that for temperature profiles typical of hydrogen-burning shells mixing of appreciable amounts of helium can only be achieved if the second Na rise is penetrated. This, however, leads to oxygen and sodium anomalies exceeding those observed (with the exception of a few stars in $\\omega$~Cen, a multi-metallicity, untypical cluster).  In terms of our normalized mass coordinate (defined such that $\\delta m = 0$ at $X = 10^{-4}$ at the bottom of the shell) this puts an observationally constrained limit for the maximum mixing depth of $\\delta m_{\\rm mix} > 0.06 \\div 0.07$.  Our complete models confirm this argument: helium enrichment in excess of $\\triangle Y_{\\rm env} \\approx 0.05$ due to deep mixing can be ruled out for those stars with Na-O-anomalies as observed in clusters such as M15, M92, M3, and even M13 which presents one of the most extended blue horizontal branch.  \\begin{figure} \\centerline{\\includegraphics[scale=0.45,draft=false]{fig13.eps}} \\caption[]{[Na/Fe] and [O/Fe] for giants in globular clusters. The symbols identify individual clusters and have the same meaning as in Fig.~\\ref{f:obsano}. Overlaid are the five theoretical predictions shown in Fig.~\\ref{f:naoanti1}, labeled by the values of $D_{\\rm mix}$. In the lower panel the two higher values relate to two lines each.} \\protect\\label{f:naomv} \\end{figure}  The details of the surface abundance history along the RGB depend on the details of the deep mixing process and thus on its nature which we did not attempt to specify here.  In particular, the field-to-cluster differences must be understood (especially the fact that field giants do not present the O-Na anticorrelation (\\cite{gsc:00}), indicating a deeper and more efficient mixing in their globular cluster counterparts) and seem to point out a non-negligible impact of environmental effects on the extra mixing efficiency.  However, we can compare the observations with the histories of Na and O abundance anomalies predicted by our simple mixing prescriptions in order to get constraints for a solid physical model.  This is done in Fig.~\\ref{f:naomv}, which shows surface abundances for several clusters as a function of brightness, i.e.\\ progressing evolution. Also shown is the theoretical prediction of the five calculations displayed in Fig.~\\ref{f:naoanti1}. Although the observational data are very few (the uncertainty in the abundances is of order 0.2 dex), and information for stars before or at the bump is available only for M13, some effects can be recognized, nevertheless.  All clusters show the whole spread of Na abundances between the canonical case without extra mixing (they would lie on a horizontal line) and that as obtained from the calculations with less deep mixing (solid lines).  As already mentioned, very low O abundances are only found in M13 giants, and the observed spread for this element is a signature of the mixing rate.  For both elements the abundance anomalies are limited to values predicted by models with extra mixing {\\em not} penetrating the second $^{23}$Na rise in the hydrogen shell. Stars with intermediate anomalies we interpret as being due to mixing penetrating less deeply into the hydrogen shell. They could be reproduced with properly adjusted mixing parameters.  Over the small brightness interval for which we have data (for all the clusters except M13 only the brightest giants are accessible), no significant abundance evolution is recognizable.  The increase in $\\rm [Na/Fe]$ at the RGB tip obtained in all calculations labeled $\\triangle X=0.20$ is not visible in the observations, which may be taken as indication that the physical reasons (e.g.\\ rotation) for the additional mixing have lost their importance or even vanished.  Such a possible time dependence of the extra mixing has not been taken into account in our simple mixing prescriptions, but should result from more physically motivated models. (We recall that in \\citen{dw:96} the carbon evolution could be reproduced with constant mixing parameters, however.) We will therefore, in a forthcoming paper, use the model by \\citen{dt:00}, which includes temporal changes in the diffusion constant due to angular momentum transport.  In our models, the largest abundance changes take place at the onset of the additional mixing, i.e., early on the RGB, after the bump (around $M_V\\approx -0.5$). This might be visible in the M13 data. The low brightness group has normal abundances, but above $M_V\\approx -0.4$ strong anomalies already appear, reminiscent of the steep increase in the calculations mixing deeper (dashed lines). In the case of Na no further enhancement is visible (the range of values does not varies along the RGB), while O seems to get depleted further. In spite of the incomplete data at hand, the abundance anomalies seem to develop rather early and within a narrow brightness range, but then do not increase any more. This general behaviour is consistent with our theoretical predictions -- though not exactly reproduced -- and again points to not too deep mixing at moderate speed. The calculations with extreme helium mixing would predict the largest O depletion and Na enhancement all along the RGB; the absence of such stars can therefore not be explained with a selection effect working against the shortest lived stars at the tip of the RGB. The fact that M13 seems to show anomalies already before the bump has to be taken with care, because we compare here only with one stellar model which has a metallicity almost a factor of 10 smaller than that appropriate for M13. At that metallicity, the bump would occur about 1~mag earlier (and our initial model is about 0.2~mag brighter than the end of the bump).  We stress the fact that the present observational data do not allow a detailed comparison with the abundance evolution, but only rough qualitative statements.  Although Fig.~\\ref{f:naomv} again demonstrates how the abundance evolution depends on the assumption about the additional mixing process, our calculations also show that the nucleosynthesis argument given above is valid for {\\em all} mixing descriptions tested and therefore model-independent. From this we predict that {\\sl red giants exhibiting the observed Na-O-anomalies do not have envelopes enriched in helium by much more than $\\triangle Y_{\\rm env} \\approx 0.05$, which is comparable to the general uncertainty in our knowledge about the helium abundances in such stars.}  Concerning the brightness reached during the RGB evolution, we showed that for specific mixing prescriptions large excess luminosities can be achieved at the end of the evolution without simultaneous mixing of large amounts of helium. As the reason we could identify the burning of the outer parts of the hydrogen shell where some H-He-mixing had happened in earlier phases. These regions become hot enough for significant hydrogen burning on such short time-scales that the deep mixing process is not able to mix the products any longer into the convective envelope. Such models experience very strong mass loss under the assumption of the continuing validity of a Reimers-type stellar wind and finish the RGB phase with total masses below $0.6\\,M_\\odot$ and envelope masses of only 10\\% of this. They could be candidates for blue HB stars and would link observed abundance anomalies, deep mixing and the second parameter problem, as suggested by \\citen{lh:95} and \\citen{swei:97}. They would avoid the problem of overproducing O-Na-anomalies as would result from the helium mixing investigated by \\citen{swei:97}.  However, we repeat our warning that the details of the evolutionary consequences of deep mixing depend crucially on the mixing process and history. Mixing speed and depth are both important for the amount of helium mixed and for the result of the competing processes ``mixing'' and ``burning'' and their time-scales. The mixing depth is also linked to the criterion for deep mixing, for which we have investigated several simple recipes.  To conclude, we need a solid physical picture for the deep mixing process in order to be able to investigate its effect on red giant evolution further. Presently, we can only point out some interesting possibilities -- such as the overluminosities -- and derive model-independent features, such as our main conclusion that the observed anomalies of oxygen and sodium rule out strong helium enhancement and therefore very deep mixing."
    },
    "0002/astro-ph0002034_arXiv.txt": {
        "abstract": "There is increasingly strong evidence that gamma-ray bursts (GRBs) are associated with star-forming galaxies, and occur near or in the star-forming regions of these galaxies.  These associations provide indirect evidence that at least the long GRBs detected by BeppoSAX are a result of the collapse of massive stars.  The recent evidence that the light curves and the spectra of the afterglows of GRB 970228 and GRB 980326 appear to contain a supernova component, in addition to a relativistic shock wave component, provide more direct clues that this is the case.  Here we establish that GRBs and their afterglows are both detectable out to very high redshifts ($z \\gtrsim 5$). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002202_arXiv.txt": {
        "abstract": "{We review recent advances in the study of binary star formation that have been made using the smoothed particle hydrodynamics technique.}  \\begin{article} ",
        "introduction": "The Smoothed Particle Hydrodynamics (SPH) numerical method was introduced by Lucy \\cite{Lucy77} and Gingold and Monaghan \\cite{GinMon77}.  Its first application was in the field of star formation, where it was used to study whether or not a rapidly-rotating polytrope could undergo fission to form a close binary system \\cite{Lucy77, GinMon78}. Since this initial application, SPH has been widely used in the study of star formation, for example \\cite{GinMon81, MiyHayNar84, Lattanzioetal85, Durisenetal86, SigKla97, BonBas91, BatBonPri95, Heller93, NelPap93, ArtLub94, Larwoodetal96, Nelsonetal98, VanCam98, Chapmanetal92}.  SPH is has many attributes which make it particularly well suited to the study of star formation. SPH is Lagrangian and does not require a computational grid. Thus, it can efficiently follow problems with large density contrasts since computational effort is not wasted simulating the low-density regions.  Also, recent SPH implementations \\cite{Evrard88, HerKat89, Benz90, Monaghan92} use spatially and temporally varying smoothing lengths so that the resolution increases automatically with increasing density; the complex multi-grid and adaptive-grid schemes that are used for finite difference methods are avoided.  In this proceedings, we review some of the recent advances in the study of binary star formation that have been made using SPH.  In Section \\ref{jeansmass}, we discuss the importance of always resolving the Jeans mass in numerical studies of self-gravitating gas.  While this has been demonstrated using various types of hydrodynamic code, we concentrate specifically on the problems that can arise if this criterion is not obeyed with SPH. In Section \\ref{collstellar}, we demonstrate that, for the first time, it is now possible to perform three-dimensional hydrodynamic calculations which follow the collapse of a molecular cloud core to stellar densities.  These calculations are performed with SPH.  Finally, in Section \\ref{accretion}, we discuss how SPH has been used to study the evolution of a protobinary system as it accretes from an infalling gaseous envelope and how this work can lead to predictions of the properties of binary stars.  \\begin{figure}[t] \\begin{center} \\epsfig{file=fig01.eps, width=13cm} \\caption{\\label{bate.forces} The ratio of the gravitational force to pressure force between two SPH particles in a Jeans-mass clump of gas of radius $R=h$.  The gravitational softening length is $\\epsilon$ and the hydrodynamic smoothing length is $h$.} \\end{center} \\end{figure} ",
        "conclusions": "The Lagrangian nature of SPH and its inherent ability to provide finer spatial resolution in regions of higher density make it a powerful tool which is ideally suited for studying star formation.  Recent advances in the study of binary star formation that have been made using SPH include: the realisation that it is essential that the Jeans mass is always resolved in numerical studies of self-gravitating gas;  the ability to perform three-dimensional hydrodynamic calculations which follow the collapse of a molecular cloud core to stellar densities; the study of the effects of accretion on a protobinary systems and the development of a code which enables the evolution of accreting binary systems to be followed $\\sim 10^6$ times faster than a full hydrodynamic calculation.  The latter of these developments has resulted in the first firm predictions of the properties of binary stars for one particular model of binary star formation."
    },
    "0002/astro-ph0002491_arXiv.txt": {
        "abstract": "We present high-resolution spectroscopy and images of a photodissociation region (PDR) in M16 obtained during commissioning of NIRSPEC on the Keck II telescope. PDRs play a significant role in regulating star formation, and M16 offers the opportunity to examine the physical processes of a PDR in detail. We simultaneously observe both the molecular and ionized phases of the PDR and resolve the spatial and kinematic differences between them. The most prominent regions of the PDR are viewed edge-on. Fluorescent emission from nearby stars is the primary excitation source, although collisions also preferentially populate the lowest vibrational levels of \\h2. Variations in density-sensitive emission line ratios demonstrate that the molecular cloud is clumpy, with an average density $n = 3\\times 10^5{\\rm \\,cm^{-3}}$. We measure the kinetic temperature of the molecular region directly and find $T_{H_2} = 930$ K. The observed density, temperature, and UV flux imply a photoelectric heating efficiency of 4\\%. In the ionized region, $n_i=5\\times10^3{\\rm\\,cm^{-3}}$ and $T_{H{\\sc ii}} = 9500$ K.  In the brightest regions of the PDR, the recombination line widths include a non-thermal component, which we attribute to viewing geometry. ",
        "introduction": "Star formation occurs in dense regions of the interstellar medium, and as a result, the environments of early-type stars often contain molecular clouds. In addition to creating an \\ion{H}{2} region, where hydrogen is predominantly ionized, stellar ultraviolet radiation creates a photodissociation region (PDR) when it interacts with molecular material.  The feedback between extant stars and the nearby clouds in PDRs is important in regulating subsequent episodes of star formation.  The UV flux determines the ionization fraction in molecular clouds, which in turn sets the ambipolar diffusion rate and thereby the rate of low-mass star formation (\\cite{McK89}).  Hollenbach \\& Tielens (1997, 1999)\\nocite{Hol97,Hol99} comprehensively review the physical and chemical processes of PDRs. Far-ultraviolet (FUV) photons ($6 < h\\nu < 13.6 {\\rm\\,eV}$) dissociate molecules and excite the gas on the surface ($A_V < 3$) of the molecular cloud. The result is spatial stratification of molecular, atomic, and ionic components of a given element. At higher densities ($n\\gs 10^4 {\\rm\\,cm^{-3}}$), collisional excitation also populates the excited states (\\cite{Ste89}).  The columns of M16 (the Eagle Nebula) are PDRs. Most of the mass in the nebula is in molecular clouds, and a variety of molecular transitions are observed directly (\\cite{Pou98}; \\cite{Whi99}).  \\h2\\ is observed on the cloud surfaces (\\cite{All99}), and at the interface with the \\ion{H}{2}\\ region, the columns are prominent in optical emission lines (\\cite{Hes96}). The young cluster NGC 6611 contains many massive stars, which illuminate the columns with total FUV flux $16 {\\rm\\,erg\\,s^{-1}\\,cm^{-2}}$ (\\cite{All99}) from a distance of about 2 pc (\\cite{Hes96}) .  We simultaneously observe the spatial and spectral relationship of the excited molecular and ionized phases of the PDR. Based on these observations (described in \\S2), we determine physical conditions in \\S3, and summarize the conclusions in \\S4. ",
        "conclusions": "} Spectroscopy of M16 with NIRSPEC offers the advantage of simultaneous observation of both the molecular and ionized phases of its photodissociation regions.  At high resolution, we detect the spatial and kinematic signatures of photoexcitation of molecular material. Most of the emission is due to distinct regions of photodissociation on the molecular cloud's surface that we observe edge-on. The spectrum includes several lines that originate in high vibrational levels, which demonstrates that fluorescent excitation predominates. The density $n\\sim10^5\\pcc$ and varies among the emitting regions, so collisions also preferentially populate low-lying vibrational levels. We measure kinetic temperature $T_{H_2}=930$ K in the molecular gas, which requires efficient photoelectric heating near the cloud's surface. In the spatially-segregated ionized region, $T_{H{\\sc ii}}=9500$ K, and the velocity widths of the brightest regions include a macroscopic contribution."
    },
    "0002/astro-ph0002172_arXiv.txt": {
        "abstract": "The temperature of the intergalactic medium (IGM) is an important factor in determining the line-widths of the absorption lines in the \\lya forest. We present a method to characterise the line-widths distribution using a decomposition of a \\lya spectrum in terms of discrete wavelets. Such wavelets form an orthogonal basis so the decomposition is unique. We demonstrate using hydrodynamic simulations that the mean and dispersion of the wavelet amplitudes is strongly correlated with both the temperature of the absorbing gas and its dependence on the gas density. Since wavelets are also localised in space, we are able to analyse the temperature distribution as a function of position along the spectrum. We illustrate how this method could be used to identify fluctuations in the IGM temperature which might result from late reionization or local effects. ",
        "introduction": "Resonant absorption by neutral hydrogen in the intergalactic medium along the line of sight to a distant quasar is responsible for the many absorption lines seen in the \\lya forest, blueward of the quasar's \\lya emission line (Bahcall \\& Salpeter 1965, Gunn \\& Peterson 1965; see Rauch 1998 for a review).  The general properties of these \\lya absorption lines are remarkably well reproduced by hydrodynamic simulations of cold dark matter (CDM) dominated cosmologies (Cen \\etal 1994, Zhang, Anninos \\& Norman 1995, Miralda-Escud\\'e \\etal 1996, Hernquist \\etal 1996, Wadsley \\& Bond 1996, Zhang \\etal 1997, Theuns \\etal 1998).  On large scales where pressure is unimportant, gas traces the dark matter and the \\lya spectrum can be used to infer the underlying density perturbations in the dark matter (Croft \\etal 1997, Nusser \\& Haehnelt 1999). On small scales however, pressure gradients oppose the infall of gas into small potential wells (Jeans smoothing), leaving the absorber more extended in space than the underlying dark matter. The width of the absorption line is then determined by residual Hubble expansion across the absorber (Hernquist \\etal 1996), Jeans smoothing and thermal broadening. Theuns, Schaye \\& Haehnelt (2000) analysed various line broadening mechanisms and demonstrated the importance of the gas temperature in controlling the line-widths.  The strong dependence of the small-scale properties of the \\lya forest on the temperature of the gas allows one to reconstruct the thermal evolution of the IGM. The gas temperature is set by the balance between adiabatic cooling caused by expansion and photo-heating by the UV-background. This introduces a tight relation between density and temperature, $T=T_0(\\rho/\\langle\\rho\\rangle)^{\\gamma-1}$ (Hui \\& Gnedin 1997). The parameters $T_0$ and $\\gamma$ of this \\lq equation of state\\rq~ are very sensitive to the reionization {\\em history} of the IGM (Haehnelt \\& Steinmetz 1998). This is because thermal time scales are long in the low density IGM probed by the \\lya forest, hence that gas retains a memory about the past history of the ionising background. Consequently, the \\lya forest provides us with a fossil record of the history of reionization, which can be explored by unravelling its thermal history as deduced from the \\lya forest.  Schaye \\etal (1999, see also Ricotti, Gnedin \\& Shull 2000) developed and tested a method to infer $T_0$ and $\\gamma$ based on the line-widths of the absorption lines. Applying this method to high resolution QSO spectra for a range of redshifts, they found (Schaye \\etal 2000) that the temperature $T_0$ decreases with decreasing redshift as expected, however, there is a large increase in $T_0$ round $z=3$, together with a decrease in the value of $\\gamma$. They attributed this change in the equation of state to late reionization of helium II. They also noted that the temperature at higher redshifts is still fairly high, which might be an indication that we are approaching the epoch of hydrogen reionization.  The method of Schaye \\etal to characterise line-widths is based on Voigt profile fitting of absorption lines (Webb 1987, Carswell \\etal 1987). The rationale behind fitting absorption lines with a Voigt profile is partly historical, and stems from earlier theoretical models in which the forest was produced by a set of \\lya \\lq clouds\\rq~. The line-width of these absorbers was assumed to be set by thermal and \\lq turbulent\\rq~ broadening, which would produce a Voigt profile, and line blending was responsible for the lines with large deviations from the Voigt profile. In the new paradigm of the \\lya forest absorption in the general IGM is responsible for lines, and there is no a priori reason to expect lines to have the Voigt shape.  In this paper we discuss a different method of characterising line-widths, based on discrete wavelets (see e.g. Press et al. 1992 for an introduction and further references). Wavelets provide an orthogonal basis for a unique decomposition of a signal (the spectrum) in terms of localised functions with a finite bandwidth. Thus they are a compromise between characterising a signal in terms of its individual pixel values and in terms of Fourier modes. In the first case, the characterisation has no information on correlations between different pixels (no frequency information) but perfect positional information. A Fourier decomposition, on the other hand, has perfect frequency information but no positional information. The analysis of a spectrum in terms of wavelets has the advantage that one can study the clustering of lines (\\lq positional information\\rq), as a function of their widths (\\lq frequency information\\rq).  The usage of wavelets to analyse QSO spectra was pioneered by Pando \\& Fang (1996, 1998), who used a wavelet analysis of \\lya absorption lines to describe the clustering of those lines. The wavelet analysis detected large scale structure in the \\lya forest, which had proved difficult using more traditional methods. In contrast to Pando \\& Fang, we will use wavelets to analyse the absorption spectrum directly, thereby eliminating the somewhat subjective step of first decomposing the continuous spectrum in absorption lines. The advantage of this new method is that it allows us to objectively characterise the typical width of absorption features as a function of position along the spectrum\\footnote{We will usually refer to absorption features as \\lq lines\\rq, but this is just a convenient name for what the eye picks out. The wavelet decomposition itself is unique and has no prejudice as to what should be considered a line.}.  We will show using hydrodynamic simulations that the probability distribution of wavelet amplitudes can be used to characterise the equation of state of the absorbing medium, in terms of the temperature at the mean density, $T_0$, and the slope, $\\gamma$, of the temperature-density relation. In addition we use the fact that wavelets are localised in position along the spectrum, thereby allowing us to detect spatial variations in $T_0$ and/or $\\gamma$, which might be present as a result of late helium II reionization or local effects.  This paper is organised as follows. In Section~\\ref{sect:setup} we first give a brief description of the generation of mock spectra from our simulations and illustrate the decomposition of the spectra in discrete wavelets. The statistics of the wavelet amplitudes for different simulations is discussed in Section~\\ref{sect:analysis} and the results are summarised in Section~\\ref{sect:conclusions}. Recently, Meiksin (2000) discussed indepently the application of wavelets to QSO spectra. ",
        "conclusions": "\\label{sect:conclusions} Clues to the thermal history of the Universe are hidden in the small scale structure of the \\lya forest. There are two reasons for this. Firstly, the widths of absorption lines are very sensitive to the temperature of gas, and secondly, thermal time scales are long in the low-density IGM that is responsible for the \\lya forest. Since the temperature of the photo-ionised IGM is determined by the evolution of the ionising background, unravelling the thermal history will have the added benefit of putting strong limits on the sources of UV light at high redshifts.  We have presented a new way of analysing the small scale structure of the \\lya forest, based on the unique decomposition of a spectrum in discrete wavelets. We have shown that the rms amplitude $\\langle A^2\\rangle$ of narrow wavelets (18.3 km $^{-1}$) correlates strongly with the temperature of the IGM, and also depends on the slope of the equation of state. We have quantified to what extent different models can be distinguished, using statistics of $\\langle A^2\\rangle$  Our mock spectra have been designed to mimick an observed spectrum of QSO 1107+485 as much as possible. In particular, we have imposed on our simulated spectra the same large scale optical depth fluctuations as are observed in QSO 1107, making our mock spectra quite realistic. Even so, we can still easily distinguish between models that differ in temperature by less than 30 per cent. We have quantified the dependence of these statistics on numerical artifacts (missing long wavelength perturbations due to the smallness of our simulation box) and on the amplitude of the dark matter fluctuations ($\\sigma_8$).  Wavelets are also localised in space, making it possible to study $T_0$ and $\\gamma$ as a function of position along the spectrum. We characterised the extent to which we can distinguish models with a single value of $T_0$ from a model with temperature fluctuations, as might result from late reionization or local effects."
    },
    "0002/astro-ph0002344_arXiv.txt": {
        "abstract": " ",
        "introduction": "By tracing the average properties of galaxies in rich clusters and their progenitors from low to high redshift, we can study how galaxies in similar environments evolve with cosmic time, and obtain insights into the phenomenon of galaxy formation. In nearby and intermediate-redshift clusters, the majority of bright galaxies are known to form a tight and conspicuous red sequence in the color--magnitude diagram (e.g., Visvanathan, Sandage 1977; Bower et al. 1992). This color--magnitude (hereafter C--M) relation seen in nearby clusters extends over a range of at least 4-magnitudes (Garilli et al. 1996) and possibly 5--6 mag (e.g., De Propris et al. 1998) from the brightest cluster galaxies; the slope of the relation does not vary from cluster to cluster.  Bower et al. (1992) showed that the small dispersion of the C--M relation for early-type galaxies in the Coma and Virgo clusters places a strong constraint on their age or coevality by assuming a simple star-formation history; they found that the early-type galaxies in these clusters are older than $\\sim$ 10 Gyr, if an age spread of $\\sim 1$ Gyr is assumed (see Bower et al. 1998 for more general discussion). Arag\\'on-Salamanca et al. (1993) then found a systematic evolution in the colors of the reddest cluster members which they interpreted as being passive evolution of old galaxies formed at $z \\gtsim 2$. Ellis et al. (1997) and Stanford et al. (1998) investigated {\\it morphologically selected} early-type galaxies in intermediate-redshift clusters, and found that the well-defined C--M sequence with small scatter holds there. Gladders et al. (1998) also found that the evolution of the slope of the C--M sequence is consistent with the predictions of passive evolution. These conclusions from studies of the C--M relation are complemented by those from Fundamental Plane analysis (e.g., van Dokkum et al. 1996, 1998). The C--M relation is more favorably interpreted as being a metallicity sequence, rather than an age sequence, since otherwise the slope should be much steeper than observed even at an intermediate redshift (Kodama, Arimoto 1997), although some age difference between bright and faint galaxies has been pointed out by some authors based on the observed differences in the absorption-line strength (e.g., Terlevich et al. 1999).  On the other hand, there are many observational indications that the star-formation activity in more distant clusters is higher than that in local clusters. It is known that the fraction of blue galaxies in clusters increases rapidly with the redshift at intermediate-redshift (Butcher, Oemler 1978, 1984; Rakos, Schombert 1995). Rakos and Schombert showed that the fraction of blue galaxies becomes $\\sim$ 80\\% at $z =$ 0.9. Through spectroscopic studies of intermediate-redshift clusters, it is also known that a significant fraction of galaxies have emission-lines and/or post-starburst signatures (e.g., Dressler, Gunn 1992; Postman et al. 1998; Dressler et al. 1999; Poggianti et al. 1999). Postman et al. (1998) investigated two clusters at $z \\sim$ 0.9 spectroscopically, and found that about half of the galaxies show high levels of star formation activity. It is expected from these results that a more active evolution of galaxies may be observed in clusters at higher redshift.  Recently, a number of clusters and cluster candidates at $z$ $\\gtsim$ 1 have been discovered (Dickinson 1995; Yamada et al. 1997; Stanford et al. 1997; Hall, Green 1998; Ben\\'\\i tez et al. 1999; Tanaka et al. 1999; Rosati et al. 1999). By observing high-redshift clusters, any differences in the star-formation history can be seen more clearly. Tanaka et al. (1999) investigated the colors of galaxies in the cluster 1335.8+2820 at $z \\sim$ 1.1, and showed that the optical and near-infrared color distributions of the galaxies in this cluster are wide, which may be attributed to variations in the star-formation activity. They found that the fraction of galaxies which show some UV excess over the expected colors for passive evolution is more than 75\\% in this cluster.  Stanford et al. (1997) investigated the NIR-selected cluster CIG J0848+4453 at $z =$ 1.27, and also showed that those galaxies whose $RJK$ colors are consistent with passively evolving elliptical models have relatively blue $B-K$ colors, which indicates that some star-formation activity is occurring in the old galaxies in the cluster core. It is therefore interesting to expand the detailed color analysis to other clusters at $z$ $\\gtsim$ 1, and to see whether these results apply generally to high-redshift clusters.  In this paper, we consider the optical and near-infrared (NIR) colors of $K'$-selected galaxies in those clusters at $z \\sim$ 1.2 near to the radio galaxy 3C 324 (Dickinson 1997a,b; Kajisawa, Yamada 1999; Kajisawa et al. 2000, hereafter as Paper I) using images obtained with the Subaru telescope and archival data from the Hubble Space Telescope (HST).  An excess of galaxy surface density in this region was recognized by Kristian et al. (1974) and by Spinrad and Djorgovski (1984), and firmly identified by Dickinson (1995). In Dickinson's (1997b) spectroscopic study, the surface-density excess was revealed to be due to the superposition of two clusters or rich groups at $z=1.15$ and $z=1.21$.  Extended X-ray emission with a luminosity comparable to that of the Coma cluster has been detected in the direction of 3C 324 (Dickinson 1997a), which suggests that at least one of the two systems is a fairly collapsed massive system. Smail and Dickinson (1995) detected a weak shear pattern in the field that may be produced by a cluster. In the following discussion, we do not distinguish the two systems at $z \\sim 1.2$, since a detailed redshift distribution of the galaxies or the relative population of them is still not available.  We thus discuss the average properties of the two clusters. Since their redshifts are close, it has little effect on our discussion about the luminosity and color distributions.  Using a $K'$-selected sample provides such advantages as (i) the selection bias at large redshift is small, and thus a direct comparison with lower-redshift clusters is possible (Arag\\'on-Salamanca et al. 1993) and (ii) the sample selection is less affected by the star-formation activity. Furthermore, the rest-frame mid-UV color provided by the HST data enables us to detect even a small amount of recent star-formation (e.g., Ellis et al. 1997; Smail et al. 1998).  We describe the observations and data reduction briefly in section 2. In section 3, we present the color-magnitude and two-color diagrams. Our conclusions and discussions are given in the last section.  In a previous Paper I, we studied the $K$-band luminosity distribution of the galaxies in these clusters near to 3C 324 using the same $K'$-band data, which is closely related to the results presented here. In summary, Paper I showed that the shape of the bright end ($K \\ltsim 20$) of the luminosity function is similar to those of lower-redshift clusters; also the measured value of the characteristic magnitude, $K^*$, follows the evolutionary trend of intermediate-redshift clusters, which is consistent with passive evolution models with $z_{\\rm form} \\gtsim 2$. At the faint end, however, the excess galaxy surface density within 40'' of 3C 324 decreases abruptly at $K \\sim 20$, which may indicate a strong luminosity segregation with distance from 3C 324 or a real deficit in the faint galaxy population in the clusters. Kajisawa and Yamada (1999) also presented results from a study of the colors and morphologies of an {\\it optically selected} sample of cluster galaxies using the same HST data. ",
        "conclusions": "We have presented the color distribution of the $K'$-band selected galaxies in the region of the clusters at $z=1.2$ near to the radio galaxy 3C 324. While the `red finger', a sequence of galaxies with $R_{\\rm F702W}-K' \\sim$ 6 and $K' \\sim$ 18--19 whose colors are consistent with ages $\\gtsim 3$ Gyr (or $z_{\\rm f} > 4.5$; see Kodama et al. 1998) was recognized in the color-magnitude diagram, we also found that this sequence is truncated at $K' \\sim 20$ mag; also, few galaxies in the `cluster' region with fainter magnitudes have similar red colors. Garilli et al. (1996) presented optical C--M diagrams ($g-r$ vs $r$) for 67 low-redshift Abell and EMSS clusters, where most of these clusters show a red sequence extending over more than 4 magnitudes. De Propris et al. (1998) presented a C--M diagram ($B-R$ vs $H$) for the Coma cluster, where the C--M relation for early type galaxies ranges over more than 5 mag. Compared with these low-redshift clusters, the `red sequence' of the 3C 324 clusters is clearly truncated.  We investigated the significance of this truncation, assuming that the red sequence population has a $K$-band luminosity distribution whose shape is identical to the cluster population as a whole. We assumed a faint-end slope of $\\alpha = -0.9$, and adopted a characteristic magnitude of $K^*=18.4$ (Paper I). The normalization was determined from the observed number of red sequence galaxies with $17 < K < 20$, and the expected number of red galaxies with $20 < K < 21$ was estimated, taking incompleteness into account. The estimated number of red galaxies with $5.7 < R_{\\rm F702W}-K' < 6.3$ and $20 < K <21$ is 3.63 in the cluster region and 4.85 for the entire region, while the observed numbers are 0 and 1, respectively.  We also investigated the $B_{\\rm F450W}-R_{\\rm F702W}$ vs $R_{\\rm F702W}-K'$ two-color diagram and compared the observed colors with models of various star-formation histories, assuming that the galaxies are indeed at $z=1.2$. We found that the large scatter in $R_{\\rm F702W}-K'$ is due to variations in both the age and star-formation activity, and that the objects with bluer $R_{\\rm F702W}-K'$ colors have younger ages than the `red finger' objects, provided that they are not dominated by foreground/background galaxies.  For moderately red objects with $4.5 < R_{\\rm F702W}-K' < 5.5$ and $B_{\\rm F450W}-R_{\\rm F702W} > 0$, the bluer $R_{\\rm F702W}-K'$ color is found to be mainly due to younger ages. There is a trend that galaxies with fainter $K'$ mag have bluer colors. Because the near-infrared $K'$ magnitude can be an approximate measure of the total stellar mass for quiescent galaxies, it may imply that the relatively low-mass galaxies are, on average, younger than the massive galaxies in the 3C 324 clusters.  The `bulge-dominated' galaxies within the cluster region seem to form a broad color-magnitude sequence whose slope is much steeper than predicted by coeval passive-evolution models. From the above discussion, this can be interpreted as an effect of the age sequence in addition to the metallicity sequence, provided that the galaxies really are cluster members. For a comparison, in figure 9, we also plot the `metallicity plus age' sequence model, using the Kodama and Arimoto code. In this model, the zero point is the same as the pure metallicity-sequence model discussed in the previous section at $R_{\\rm F702W}-K'$ = 6 ($z_{\\rm f} = 4.5$, $\\sim 3$ Gyr old), and the fainter galaxies are 0.5 Gyr younger per 1 $M_{V}$-magnitude (at $z=0.02$). The mean metallicity at each magnitude is adjusted so that the model sequence coincides with the Coma cluster C--M relation, when it is evolved passively to $z = 0.02$. It can be seen that this model well reproduces the sequence of the `bulge-dominated' cluster-region galaxies. The age difference may thus be as much as $\\sim 2$ Gyr between galaxies with $K'=18$ and $K'=21$.  Is this picture consistent with the well-defined C--M sequence observed in intermediate-redshift clusters? The C--M relation of morphologically selected early-type galaxies in clusters at $z =$ 0.2--0.8 is known to extend over a range of at least 3--4 magnitudes with a scatter smaller than 0.1 mag (e.g., Ellis et al. 1997; Stanford et al. 1998). This is very different from the C--M sequence observed at $z=1.2$.  However, if the relatively large scatter and steep slope of the C--M relation for the bulge-dominated galaxies in the 3C 324 clusters are caused by an age difference and the evolution is passive after $z=1.2$, the younger (bluer) galaxies would change their color and luminosity more rapidly than the old (red) galaxies, and may become as red as and even fainter within a relatively short time. The C--M relation would then become tighter and would extend over a larger magnitude range, with a shallower slope. Note that the effect of an additional small amount of star formation at $z = 1.2$ with a standard IMF is negligible at intermediate-z, if the star formation ceases immediately after $z =$ 1.2, because the brightness of these component declines rapidly as massive stars die and the main sequence burns down.  For example, we compared the color and luminosity evolution of the 1 Gyr burst models with different ages, using the GISSEL96 code. One model is 3 Gyr old at $z = 1.2$ (formation redshift $z_{\\rm f} \\sim 5$, $H_{0} =$ 50 km s$^{-1}$ Mpc$^{-1}$, $q_{0}$ = 0.5) and its optical--NIR color is $R_{\\rm F702W} - K' \\sim 6$, which roughly corresponds to the `red sequence' population. The other is 1.5 Gyr old at $z = 1.2$ ($z_{\\rm f} \\sim 2$) and its $R_{\\rm F702W} - K' \\sim 5$, which represents the bluer galaxies. At $z \\sim 0.7$, the difference in $R_{\\rm F702W} - K'$ color between these two models is only $\\Delta (R_{\\rm F702W}-K') \\sim 0.1$, and the difference in $B_{\\rm F450W} - R_{\\rm F702W}$ color, which is more sensitive to a small amount of star formation, becomes $\\sim 0.1 $ by $z \\sim$ 0.5. The fading of the younger model galaxy in the $K'$ band is about 0.5 mag greater than that of the old model galaxy at $z \\ltsim$ 0.9; if we assume that the younger galaxy is about 1 magnitude fainter at $z = 1.2$, it becomes about 1.5 magnitude fainter at an intermediate redshift. From these results, the bluer `bulge-dominated' population in the 3C 324 clusters is expected to form a fainter part of the tight C--M relation at intermediate-redshift, which may at least partially explain the observed apparent truncation of the red sequence at $z=1.2$.  In Paper I, we showed that the excess of the surface density in the cluster region drops abruptly at $K \\sim$ 20 mag. One possible interpretation is that there may be an intrinsic deficit of faint galaxies in the clusters. Despite the existence of color selected galaxies with $K > 20$ mag in the cluster region, this possibility cannot be completely ruled out, since the number of these objects may not be large enough (subsection 3.2). There is, however, another possible explanation, that there is a difference in the spatial distribution between bright and faint galaxies, and that the apparent deficit of the surface-density excess is due to the more uniform distribution of the faint cluster galaxies. Indeed, such luminosity segregation is also seen in the spatial distribution of the color-selected galaxies that may be at $z \\gtsim 1.2$, as shown in figure 10 (subsection 3.3). This tendency is insensitive to subtle changes of the boundary between the `$z=1.2$' and `foreground' color domains in figure 4. There are several possible origins for such a luminosity segregation. Within the framework of the theory of biased galaxy formation with cold dark matter, luminosity segregation is naturally predicted (Valls-Gabaud et al. 1989). In biased CDM theory, the most massive and oldest galaxies correspond to the highest peaks of the initial density fluctuations, and are concentrated more strongly at the early formation epoch. Alternatively, merger events may play important roles (e.g., Fusco-Femiano, Menci 1998); since galaxy merging may occur more frequently at the dense cluster core, massive galaxies can be formed efficiently in the central region.  The difference in the sky distribution between those galaxies with different colors and magnitudes may be the result of the superposition of two systems, namely those at $z =$ 1.21 and 1.15, that may contain different galaxy populations. Postman et al. (1998) indeed show such an example.  The cluster Cl 0023+0423 at $z =$ 0.84 was revealed to consist of two poor clusters or groups of galaxies; they found that there are few luminous galaxies in one system, while luminous red galaxies exist in the other system. The 3C 324 clusters may be in similar situation to Cl 0023+0423. Although they are not associated with each other in the real space in the case of the 3C324 clusters, there may be a large variety of galaxy populations among the high-redshift clusters or groups.  Finally, we note a rather speculative feature about the galaxy spatial distribution. The large-scale distribution of those galaxies with a `$z =$ 1.2' color domain in the field extends in an east--west direction centered at 3C 324, and is aligned with the radio axis of 3C 324. This may imply a possible relationship between the past radio jet activity of 3C 324 and galaxy formation (e.g., West 1994).\\\\ \\vspace{0.5cm}  We wish to thank Alfonso Arag\\'on-Salamanca, the referee of this paper, for his invaluable comments. The present result is indebted to all the members of the Subaru Observatory, NAOJ, Japan. We thank Nobuo Arimoto for kindly providing the Kodama and Arimoto evolutionary synthesis models. This research was supported by grants-in-aid for scientific research of the Ministry of Education, Science, Sports and Culture (08740181, 09740168). This work was also supported by the Foundation for the Promotion of Astronomy. This work is based in part on observations with the NASA/ESA Hubble Space Telescope, obtained from the data archive at the Space Telescope Science Institute, which is operated by AURA, Inc.\\ under NASA contract NAS5--26555. The Image Reduction and Analysis Facility (IRAF) used in this paper is distributed by National Optical Astronomy Observatories, operated by the Association of Universities for Research in Astronomy, Inc., under contact to the National Science Foundation."
    },
    "0002/astro-ph0002422_arXiv.txt": {
        "abstract": "Using a high-resolution cosmological hydrodynamic simulation, we present a method to constrain extinction due to intergalactic grey dust based on the observed magnitudes of distant Type IA supernovae. We apply several simple prescriptions to relate the intergalactic dust density to the gas density in the simulation, thereby obtaining dust extinctions that may be directly compared to the observed distribution of supernova magnitudes.  Our analysis is sensitive to the spatial distribution of grey dust, but is not dependent on its intrinsic properties such as its opacity or grain size.  We present an application of our technique to the supernova data of Perlmutter et al., who find that their high redshift sample is $\\sim 0.2$ magnitudes fainter than the expectation for a non-accelerating, low-density universe. We find that for grey dust to be responsible, it must be distributed quite smoothly, e.g., tracing intergalactic gas. More realistic dust distributions, such as dust tracing the metal density, are inconsistent with observations at the $1.5-2 \\sigma$ level.  Upcoming observations and improved modelling of the dust distribution should lead to stronger constraints on intergalactic grey dust extinction. ",
        "introduction": "Recent observations of Type IA Supernovae (SNe) at redshifts up to $z\\sim 0.8$ (\\cite{ri98}; Perlmutter et al. 1999, hereafter \\cite{per99}) have made possible classical cosmological tests that require standard candles, such as the magnitude-redshift relation.  The most dramatic result is that these SNe appear dimmer (by $\\sim 0.2$ magnitudes) at high redshift than would be predicted in a non-accelerating universe, suggesting at face value that we live in an accelerating universe.  However, other explanations are possible, including the one we consider here, namely that distant SNe are dimmer due to extinction by intergalactic dust.  As distant standard candles, SNe are sensitive probes of extinction in the intergalactic medium (IGM).  The distribution of matter in the IGM may now be modelled accurately in the context of modern cosmology using hydrodynamic simulations (see e.g., \\cite{cen94}; \\cite{her96}; \\cite{dav99}).  The resulting IGM is not smooth, but rather traces large-scale structure.  If such structure contains not only dark matter and gas but also dust, this would result in significant variations in SNe brightnesses due to intervening extinction.  Observationally, the distribution of Type Ia SNe magnitudes has a very small dispersion (\\cite{per99}).  Thus by comparing simulations to the distribution of observed brightnesses, we can set limits on the amount of dust extinction and possibly constrain its spatial distribution with respect to intergalactic gas.  In this {\\it Letter} we present a technique for doing this, and apply it to the SNe observations of \\cite{per99}.  Intergalactic grey dust has been examined in a series of papers by Aguirre (1999a,b; hereafter \\cite{agu99}) and \\cite{agu99c}, who develop a scenario in which small grains are preferentially destroyed during ejection from galaxies, polluting the IGM with large dust grains that are effectively grey in the bandpasses of the SNe data.  This greyness is necessary in order not to violate tight limits on reddening from SNe data, which imply that galactic-type dust would provide negligible absorption (\\cite{per99}).  Furthermore, a significant fraction of the dust must reside in the IGM.  If the grey dust causing extinction were present only in the ISM of the supernova host galaxy, this would introduce too large a dispersion in observed SNe magnitudes (\\cite{ri98}).  In this study, we assume that grey dust blends smoothly from galaxies into the surrounding IGM.  Our analysis is insensitive to the intrinsic properties of the dust, such as its opacity and grain size, since we use the simulations to directly translate the observed SNe magnitude distribution into a dust extinction in magnitudes.  It is, however, sensitive to the way in which dust traces the distribution of gas in the IGM, and we will consider several simple but plausible variations of this relation.  In \\S~\\ref{sec: sims} we describe our simulations of the IGM, and of dust extinction.  In \\S~\\ref{sec: anal} we describe our analysis method and results, including constraints on grey dust afforded by current SN observations.  In \\S~\\ref{sec: disc} we discuss systematic uncertainties, and the implications of our results. ",
        "conclusions": "\\label{sec: disc}  We have found that current SNe datasets appear to have some power to constrain grey dust models.  We find it somewhat unlikely (at the $1.5-2$ level) that sufficient grey dust could be distributed in the relatively clumpy fashion expected for the metal distribution (c.f., \\cite{gne97}, \\cite{dav98}, \\cite{cen99}) to reconcile the \\cite{per99} data with a non-accelerating universe.  This suggests that any substantial grey dust component must be largely segregated from the galaxies where it was formed.  In the specific grey dust model of \\cite{agu99}, sputtering does destroy dust more effectively in denser regions, but once in the IGM the large grains have long lifetimes, so that dust is still likely to trace the metals. The question of how the dust and metals are distributed can be answered self-consistently by modelling the relevant physical processes (dust and metal ejection from galaxies) directly in the simulations (e.g., Aguirre et al., in preparation).   There are many alternative explanations for SNe appearing dimmer in the past.  For example, there may be metallicity effects in the host galaxy (\\cite{hof99}), intrinsic evolution (\\cite{rie99}), observational selection differences between CCDs used for distant SNe and photographic plates used for nearby samples (\\cite{how99}), or time evolution of the gravitational constant (\\cite{ame99}; \\cite{gar99}). Such effects would alter the interpretation of our parameter $M_{C}$. Gravitational lensing magnification (\\cite{met99}) would also change the magnitude dispersion at higher redshifts.   The study we have presented is reasonably general, in that our analysis is not dependent on the (unknown) microscopic properties of this hypothetical intergalactic dust.  Still, there are some possible systematic uncertainties, which we now consider.  We assume (as do \\cite{per99}) that the observed dispersion in SNe magnitudes in excess of the estimated observational errors is an intrinsic property of SNe and does not vary with redshift.  One could envision scenarios in which the intrinsic dispersion is lower at high redshift, thus allowing more dispersion from dust. However, as mentioned earlier, most of the statistical power of our analysis comes from the {\\it shape} of the distributions, which is not significantly affected by a change in the intrinsic dispersion. Also, we find that when we decrease the intrinsic SN dispersion ($\\sigma_{\\rm int}$), models with more dust fit slightly better at high-$z$, but the low-$z$ fit becomes worse, a trade-off which means that the overall results hardly change. The distribution shape depends on our assumption that the intrinsic dispersion has a Gaussian distribution in magnitudes, whereas a distribution skewed to fainter magnitudes could weaken our constraints. With a larger sample of low-z SNe we might be in a position to test this, by using the distribution of low-z SN magnitudes to make a simulated high-z sample with the correct intrinsic distribution shape. Also, we have made a simplifying approximation in our simulated high-z samples, by using dust extinction magnitudes that result from integrating the dust contribution from $z=0$ to exactly $z=0.5$. The real SNe are at a range of redshifts; our constraints are effectively conservative because of this.  We have also assumed that the total mass of dust does not change with redshift from $z=0.5$ to $z=0$; dust increasing with time would strengthen our constraints, while a decrease seems implausible.  Finally, the truncation of our simulated magnitude distribution at $+1$ mag is a rather approximate procedure. We have tried changing the cutoff, and find that with none (an unrealistic case) the constraints become stronger, as we might expect. With a lower cutoff, the right side of the observed distribution is not reproduced. One additional degree of freedom would involve changing the functional form of the cutoff. In the future, as observations improve, we plan to simulate such observational selection effects more carefully.  On the simulation side, our relatively low mass resolution ($m_{\\rm gas}=8.5 \\times 10^8 M_{\\odot}$) means that we miss fluctuations in the extinction that occur on smaller mass scales.  If we had higher resolution, this should have the effect of making our dust constraints tighter. We find that the dispersion in projected magnitudes largely depends on {\\it small} scale fluctuations.  We tested this by splitting the simulation into small sub-volumes and then shuffling them to remove large scale correlations (see e.g., \\cite{bl88}) before projection, and we obtained similar results.  Changing the assumed cosmological model could have some effects, although it is difficult to conceive of models which have less power on small scales while fitting other constraints (see e.g., \\cite{wc2000}).  The models of \\cite{agu99} have grains that do produce some reddening, typically in the infrared.  A complementary approach to ours is therefore to use information from different color bands. Such an approach was used by \\cite{per99}, for example, to constrain normal galactic-type dust.  Recent near-IR observations by \\cite{ri00} have made reddening constraints tight even for non-standard large dust grains.  Our approach does not make any use of color information, and so constrains the most extreme scenario in which the dust is totally ``grey''.  In summary, we have used cosmological hydrodynamic simulations to explore how intergalactic grey dust could affect observations of high redshift supernovae, and how supernova data can constrain grey dust extinction.  We conclude that only a fairly smooth distribution of dust could readily mimic the effect of an accelerating Universe.  Such a distribution would be strange, as the dust would be strongly biased away from metal producing regions.  More realistic dust distributions are mildly disfavored, but upcoming samples of SNe (at current redshifts) should enable us to put tighter limits as they more precisely determining the shape of the brightness distribution. At higher redshifts ($z\\ga 1$), grey dust predicts that SNe should show increased dimming, while decreased dimming would be strong evidence for a cosmological constant.  \\bigskip"
    },
    "0002/astro-ph0002108_arXiv.txt": {
        "abstract": "We discuss new {\\it Hubble Space Telescope} spectroscopy of the radio-loud quasar, 3C~288.1. The data cover $\\sim$590 \\AA\\ to $\\sim$1610 \\AA\\ in the quasar rest frame. They reveal a wealth of associated absorption lines (AALs) with no accompanying Lyman-limit absorption. The metallic AALs range in ionization from \\ion{C}{3} and \\ion{N}{3} to \\ion{Ne}{8} and \\ion{Mg}{10}. We use these data and photoionization models to derive the following properties of the AAL gas: 1) There are multiple ionization zones within the AAL region, spanning a factor of at least $\\sim$50 in ionization parameter. 2) The overall ionization is consistent with the ``warm'' X-ray continuum absorbers measured in Seyfert 1 nuclei and other QSOs. However, 3) the column densities implied by the AALs in 3C~288.1 are too low to produce significant bound-free absorption at any UV--X-ray wavelengths. Substantial X-ray absorption would require yet another zone, having a much higher ionization or a much lower velocity dispersion than the main AAL region. 4) The total hydrogen column density in the AAL gas is $\\log N_{\\rm H} ({\\rm cm}^{-2})\\approx 20.2$. 5) The metallicity is roughly half solar. 6) The AALs have deconvolved widths of $\\sim$900~\\kms\\ and their centroids are consistent with no shift from the quasar systemic velocity (conservatively within $\\pm$1000 \\kms ). 7) There are no direct indicators of the absorber's location in our data, but the high ionization and high metallicity both suggest a close physical relationship to the quasar/host galaxy environment.  Finally, the UV continuum shape gives no indication of a ``blue bump'' at higher energies. There is a distinct break of unknown origin at $\\sim$1030 \\AA , and the decline toward higher energies (with spectral index $\\alpha\\approx -1.73$, for $f_{\\nu}\\propto \\nu^{\\alpha}$) is even steeper than a single power-law interpolation from 1030 \\AA\\ to soft X-rays. ",
        "introduction": "Associated absorption lines (AALs) in quasar spectra provide unique information on the kinematics, physical conditions and elemental abundances in the gas near quasars. AALs are defined empirically as having 1) relatively narrow profiles (less than a few hundred \\kms ), and 2) absorption redshifts, $z_a$, near the emission redshift, $z_e$, (generally within 3000--5000~\\kms , see \\cite{wey79,fol86,fol88}). The first criterion distinguishes AALs from the class of broad absorption lines (BALs), which have velocity widths and maximum displacements that often exceed 10,000~\\kms . BALs clearly form in high-velocity winds from the central engines (see the reviews by \\cite{wey97,wey95,tur95,wey85}). AALs (or \\zaz\\ systems) can form potentially in a variety of locations --- from outflows near the black hole/accretion disk, perhaps like the BALs, to intervening gas or galaxies at large (cosmologically significant) distances. A few AAL systems are known to form in quasar ejecta (probably within a few pc of the energy source, \\cite{h97a,bar97a,gan99}), while others clearly probe extended regions on host-galaxy scales ($>$1 kpc distant, \\cite{wil75,sar82,mor86,tri96,bar97b}). Surprisingly little else is known about the nature of the absorbing regions or their relationship to other quasar phenomena.  One important clue is that both AALs and BALs in the UV appear to correlate with the presence of continuous absorption in soft X-rays (\\cite{gre96,cre99,mat98,mat99,gal99,bra99} and references therein). The X-ray absorbers in AAL sources (both quasars and Seyfert 1 galaxies) tend to have high total hydrogen column densities ($\\log N_{\\rm H} ({\\rm cm}^{-2})\\sim 21$ to 23) and high ionizations (being dominated by absorption edges of O~{\\sc vii} and O~{\\sc viii} near 0.8 keV) (\\cite{rey97,geo98a}). In contrast, the AALs typically indicate lower column densities and lower levels of ionization (cf. \\cite{h97}). The X-ray absorbers that accompany BALs have even larger total column densities of $\\log N_{\\rm H} ({\\rm cm}^{-2})\\ga 23$ (\\cite{gre96,gal99}). In those objects, the total column densities derived from X-rays exceed prior estimates from the BALs by 2 or more orders of magnitude (\\cite{h98a}). Clearly, we must consider the UV and X-ray data together to obtain a complete census and understanding of the absorbing environments.  A key question now is the physical relationship between the UV and X-ray absorbers. Mathur \\etal (1998 and refs. therein) argue that they could reasonably identify a single absorbing medium, while other work has shown that multiple regions (having different velocities, ionizations and/or column densities) are at least sometimes present (e.g. \\cite{kri96,h97a,rey97,geo98b,mat99}). High-ionization UV lines, such as \\ion{Ne}{8} \\lam\\lam 770,780 and \\ion{Mg}{10} \\lam\\lam 610,625, can directly test the UV--X-ray relationship because their ionization requirements are similar to the \\ion{O}{7} and \\ion{O}{8} edges measured in X-rays. The far-UV spectra needed to reach the \\ion{Ne}{8} and \\ion{Mg}{10} lines also encompass many under-utilized diagnostics such as the \\ion{H}{1} Lyman limit, the Lyman series lines, and numerous metal lines spanning a wide range of ionizations. Unfortunately, these features are difficult or impossible to measure in many sources because of their short wavelengths. For example, they are obscured by Galactic Lyman-limit absorption in low-redshift sources (e.g. in all Seyfert galaxies) and contaminated by the dense ``forest'' of \\Lya\\ absorbers at high redshifts. Another problem at high redshifts is that the dominant \\ion{O}{7} and \\ion{O}{8} edges are shifted out of the sensitive energy range of current X-ray telescopes.  Intermediate-redshift quasars ($0.5\\la z\\la 1$) provide a unique opportunity to measure all of the key UV and X-ray features in the same object. We have begun a program to obtain UV and X-ray spectra of several such objects. Here we discuss new UV observations of the radio-loud quasar, 3C~288.1. This source has both the required moderate redshift, $z_e\\approx 0.961$ (\\cite{sch68}), and a strong AAL system (\\cite{will95}). It also has a dramatically bipolar (lobe-dominated) radio morphology (\\cite{rei95,aku94}). The X-ray absorption properties of 3C~288.1 are not yet known, but it was weakly detected in soft X-rays with the {\\it Einstein} IPC (from $\\sim$0.16 to $\\sim$3.5~keV, \\cite{zam81}). Wilkes \\etal (1994) estimate its 2-point power-law index between 2500 \\AA\\ and 2~keV to be $\\alpha_{ox} = -1.51$ (where $f_{\\nu}\\propto \\nu^{\\alpha_{ox}}$). ",
        "conclusions": "Comparative UV and X-ray studies are beginning to yield new insights into the nature of AGN absorbing environments. However, the the physical relationship of the UV and X-ray absorbers remains unclear. 3C~288.1 provides another counterexample to simple models that would attribute all of the UV and X-ray features to identically the same gas (see \\S1 for references). Realistic models must include the increasing levels of complexity implied by new and better data. Intermediate-redshift quasars will be an important source of new data because they allow us to measure a wide range of under-utilized features in the rest-frame far UV --- for example  the \\ion{H}{1} Lyman limit and high-ionization AALs like \\ion{Ne}{8} \\lam\\lam 770,780 and \\ion{Mg}{10} \\lam\\lam 610,625. Our analysis of these features in 3C~288.1 shows that multiple absorbing regions (spanning a range of ionizations, densities and/or distances from the quasar) contribute to the AAL spectrum (\\S4.3). Moreover, the column densities inferred from the AALs are too small to produce significant bound-free opacity at any UV through X-ray wavelengths (\\S4.5). There will be no significant X-ray absorption in this object unless it occurs in yet another region --- having a much higher ionization or a much lower velocity dispersion than the dominant AAL absorber. We are now pursuing X-ray observations of 3C~288.1 to test these predictions.  One essential ingredient for any model of these regions is the location of the absorbing gas. Indirect evidence suggests that the AALs in 3C~288.1 form close to the quasar, probably well within the radius of the quasar's host galaxy (\\S5.2). Unfortunately, that weak constraint on the location allows for a variety of possible absorption sites, e.g. interstellar gas in the extended host galaxy, a dense torus surrounding the active nucleus, or an outflow from the central engine/accretion disk. The AAL kinematics indicate that the absorber is clearly {\\it not} part of a high-velocity wind like the BALs; the lines are only $\\sim$900 \\kms\\ wide and their centroids are within $\\pm$1000 \\kms\\ of the quasar rest frame (\\S4.1). Better constraints on the location and kinematics will require higher resolution spectra and, ideally, repeated observations at both UV and X-ray wavelengths (to test for variable absorption, cf. \\cite{h97a,bar97a,geo98b}). Those data would also provide more complete and more reliable estimates of the ionizations, kinematics, abundances and column densities in the different absorbing zones.   \\bigskip"
    },
    "0002/astro-ph0002087_arXiv.txt": {
        "abstract": "The properties of old globular cluster systems in galaxy halos are used to infer quantitative constraints on aspects of star formation that are arguably as relevant in a present-day context as they were during the protogalactic epoch. First, the spatial distribution of globulars in three large galaxies, together with trends in total cluster population vs.~galaxy luminosity for 97 early-type systems plus the halo of the Milky Way, imply that bound stellar clusters formed with an essentially universal efficiency throughout early protogalaxies: by mass, always $\\epsilon_{\\rm cl}=0.26\\%\\pm0.05\\%$ of star-forming gas was converted into globulars rather than halo field stars. That this fraction is so robust in the face of extreme variations in local and global galaxy environment suggests that any parcel of gas needs primarily to exceed a {\\it relative} density threshold in order to form a bound cluster of stars. Second, it is shown that a strict scaling between total binding energy, luminosity, and Galactocentric position, $E_b=7.2\\times10^{39} \\,{\\rm erg}\\,(L/L_\\odot)^{2.05}\\,(r_{\\rm gc}/8\\,{\\rm kpc})^{-0.4}$, is a defining equation for a fundamental plane of Galactic globular clusters. The characteristics of this plane, which subsumes {\\it all other} observable correlations between the structural parameters of globulars, provide a small but complete set of facts that must be explained by theories of cluster formation and evolution in the Milky Way. It is suggested that the $E_b(L,r_{\\rm gc})$ relation specifically resulted from star formation efficiencies having been systematically higher inside more massive protoglobular gas clumps. ",
        "introduction": "\\label{sec:1}  The old globular cluster systems (GCSs) in galaxy halos are of especial interest in a conference such as this, as they stand potentially to shed light on both local and global star formation in both protogalactic and present-day settings. Because of their great age, their ubiquity in galaxies of most any Hubble type, and their impressive spatial extent (globulars can be found as far as 100 kpc away from the centers of large galaxies), the integrated properties of GCSs clearly had to have been influenced by very general, large-scale star-formation processes in the early universe. But at the same time, most new stars today, whether in the Galactic disk or in galaxy mergers and starbursts, are born not in isolation but in groups. To be sure, the formation of a bona fide star cluster is a {\\it rare} event (only a very small fraction of young stellar groups emerge from their natal clumps of gas as gravitationally bound units) but, as will be argued here, it is one that occurs {\\it regularly}. Thus, individual globular clusters must also be viewed---even if in a limit---as the products of a robust mode of smaller-scale star formation that has always been viable.  It should come as no surprise, then, that there is no definitive theoretical description of the formation history of globular clusters; while some concepts have been identified that do seem likely to survive as elements of a correct theory in the future, at this point there is simply no model that can claim completeness. Detailed discussions of the many theories already in the literature may be found in, e.g., \\cite*{ash98} or \\cite*{mey97}. The focus here will instead be on recent progress in extracting quantitative and (as nearly as possible) model-independent constraints from the data on many GCSs. A similarly empirical discussion is given by \\cite*{har00}, with an eye mostly to implications for large-scale aspects of galaxy formation and evolution (see also \\cite{ash98}). In what follows, particular emphasis will be placed on applications to the problem of {\\it generic} star formation on subgalactic scales.  A serious concern, when trying to use GCSs in this way, is the influence of dynamical evolution on the gross properties of cluster systems that have been immersed for a Hubble time in the tidal fields of their parent galaxies. Two-body relaxation and evaporation, disk- and bulge-shock heating, chaotic scattering or disruption by a compact nucleus, and dynamical friction: all of these processes whittle away at clusters individually and collectively (\\cite{spi87}; \\cite{agu88}; \\cite{ost89}; \\cite{cap93}; \\cite{mwa97}; \\cite{mwb97}; \\cite{mwc97}; \\cite{ves97}). In the case of the Milky Way particularly, the net effect is to define a roughly triangular region in mass--radius space, within which globulars are predicted to survive a 10-Gyr dynamical onslaught (\\cite{fal77})---the implication being that any clusters born outside such a ``survival triangle'' would have disappeared, taking with them vital information on their birth properties. However, further investigation shows that Milky Way clusters located at large Galactocentric radii ($r_{\\rm gc}\\ga 3$ kpc, roughly the effective radius of the bulge) {\\it do not fill} their expected survival triangles (\\cite{cap84}; \\cite{gne97}). Similarly, evolutionary calculations geared to conditions in the giant elliptical M87 (\\cite{mwb97}) suggest that the damage done to an initial GCS may be largely confined to within an effective radius in that galaxy as well. Thus, there are some features of globular clusters, and of GCSs, that are {\\it not} dictated purely by evolution from some completely obscured initial conditions, and while care must be taken in identifying such observables, the task is not an impossible one.  Following the suggestion that large fractions of the GCSs in ellipticals may have formed in major mergers (e.g., \\cite{sch87}; \\cite{ash92}), and the related discovery of young, massive star clusters in systems like the Antennae galaxies (\\cite{whi95}), much recent discussion in this field has centered on the interpretation of (often) bi- and multimodal distributions of globular cluster colors (as indicators of integrated metallicity) in relation to their host galaxies' dynamical histories (e.g., \\cite{zep93}; \\cite{for97}; \\cite{cot98}; \\cite{kis98}; \\cite{kig99}; \\cite{cot00}). However, as will be discussed further in \\S\\ref{sec:3}, metallicity is completely decoupled from the other basic properties of individual globulars in the Milky Way; thus, while GCS colors are of interest in the context of galaxy formation and chemical evolution, they seem to be of marginal importance to the star-formation process itself. In this connection, it is worth noting explicitly that literally thousands of genuinely old globulars have now been identified, in many galaxies besides our own, with colors as red as or redder than solar: {\\it high metallicity has never posed any apparent obstacle to the formation of massive star clusters}. This cautions against theories of globular formation such as those of \\cite*{fal85} and \\cite*{mur93}, which, at least in their current form, are able to account only for metal-poor objects.  Of rather more direct relevance to the small-scale problem are the luminosity or mass functions of cluster systems. It is well known that the overall range of Galactic globular cluster masses, $m\\sim 10^4$--$10^6\\ M_\\odot$, and the mean value $\\langle m\\rangle=2.4\\times10^5\\ M_\\odot$, are essentially universal properties of other GCSs. Traditionally, theories of globular cluster formation attempted to explain just this basic mass scale (e.g., \\cite{pee68}; \\cite{fal85}). More recently, however, attention has turned to the full {\\it mass spectrum}, $d{\\cal N}/dm$ (the number of clusters with mass $m$ in a single galaxy), which---in the regime usually observed, $m\\ga 10^5\\ M_\\odot$---is similar (though not identical) from galaxy to galaxy; shows no detectable variations with radius in any one system; and is remarkably similar to the mass function, $d{\\cal N}/dm\\propto m^{-1.7}$ or so, of the dense gas clumps currently forming stars in the giant molecular clouds of the Milky Way disk (\\cite{har94}). These facts have suggested a general physical picture, developed by \\cite*{har94} building on earlier arguments by \\cite*{lar88} and \\cite*{lar93}, in which handfuls of globulars form in each of many protogalactic fragments whose gas masses (of order $10^9\\ M_\\odot$), sizes, and internal velocity dispersions correspond to disk GMCs scaled up by some three orders of magnitude in mass. Such a picture has the obvious appeal of being at least conceptually consistent both with current models of hierarchical galaxy formation (there is clearly some affinity with the classic picture of \\cite{sea78} as well) and with the observed pattern of present-day star formation. It also has more specific attractions (e.g., the protogalactic fragments and their protoglobulars are presumed to be supported largely by nonthermal mechanisms, thus allowing for metal-rich protoglobulars with the correct mass scale) but many details remain to be worked out. See \\cite*{mcl96} (and compare the rather different view of \\cite{elm97}) for an attempt at a quantitative theory for GCS mass spectra which, with the \\cite*{har94} framework as a backdrop, makes explicit use of the emerging links between present-day and protogalactic star formation.  In addition to color and luminosity distributions, photometric studies of extragalactic GCSs yield estimates of their specific frequencies (total cluster populations, normalized to the parent galaxy luminosity) and spatial structures (number density of globulars as a function of galactocentric position). These are the focus of \\S\\ref{sec:2}, where it is shown that a large sample of early-type galaxies display a common ratio of GCS mass to total mass in stars and (hot) gas---an apparently universal efficiency for the formation of globular clusters (\\cite{mcl99}). Finally, \\S\\ref{sec:3} discusses the binding energies of globular clusters in the Milky Way (\\cite{mcl00}). It is shown that a strong correlation between binding energy, total cluster luminosity, and Galactocentric position is instrumental in defining a fundamental plane for Galactic globulars---which, by incorporating every one of their many other structural correlations, systematically reduces these data to the smallest possible set of {\\it independent} physical constraints for theories of cluster formation and evolution. ",
        "conclusions": ""
    },
    "0002/astro-ph0002278_arXiv.txt": {
        "abstract": "Gamma-ray bursts, if they are generated in the process of interaction between relativistic strongly magnetized winds and an ambient medium, may be accompanied by very short pulses of low-frequency radio emission. The bulk of this emission is expected to be at the frequencies of $\\sim 0.1-1$ MHz and cannot be observed. However, the high-frequency tail of the low-frequency radio emission may reach a few ten MHz and be detected, especially if the strength of the magnetic field of the wind is extremely high. ",
        "introduction": "The prompt localization of gamma-ray bursts (GRBs) by BeppoSAX led to the discovery of X-ray/optical/radio afterglows and associated host galaxies. Subsequent detections of absorption and emission features at high redshifts ($0.43\\leq z \\leq 3.42$) in optical afterglows of GRB and their host galaxies clearly demonstrate that at least some of the GRB sources lie at cosmological distances (for reviews, see \\cite{Piran99}; \\cite{Vietri99}).  A common feature of all acceptable models of cosmological $\\gamma$-ray bursters is that a relativistic wind is a source of GRB radiation. The Lorentz factor, $\\Gamma_0$, of such a wind is about $10^2-10^3$ or even more (e.g., \\cite{Fenimore93}; \\cite{Harding97}). A very strong magnetic field may be in the plasma outflowing from cosmological $\\gamma$-ray bursters (\\cite{Usov94}; \\cite{Blackman96}; \\cite{Vietri96}; \\cite{Katz97}; \\cite{Meszaros97}). A relativistic, strongly magnetized wind interacts with an ambient medium (e.g. an ordinary interstellar gas) and decelerates. It was pointed out (\\cite{Meszaros92}) that such an interaction, assumed to be shock-like, may be responsible for generation of cosmological GRBs.  The interaction between a relativistic, strongly magnetized wind and an ambient plasma was studied numerically by Smolsky \\& Usov (1996, 2000) and Usov \\& Smolsky (1998). These studies showed that about 70\\% of the wind energy is transferred to the ambient plasma protons that are reflected from the wind front. The other $\\sim 30$\\% of the wind energy losses is distributed between high-energy electrons and low-frequency electromagnetic waves that are generated at the wind front because of nonstationarity of the wind---ambient plasma interaction (see below). High-energy electrons, accelerated at the wind front and injected into the region ahead of the front, generate synchro-Compton radiation in the fields of the low-frequency waves. This radiation closely resembles synchrotron radiation and can reproduce  the non-thermal radiation of GRBs observed in the Ginga and BATSE ranges (from a few KeV to a few MeV).  \\cite{Ginzburg73} and \\cite{Palmer93} suggested that GRB might be sources of radio emission, and that it might be used to determine their distances and, through the dispersion measure, the density of the intergalactic plasma. Here we consider some properties of the low-frequency waves generated at the wind front.  We argue that coherent emission by the high-frequency tail of these waves may be detected, in addition to the high-frequency (X-ray and $\\gamma$-ray) emission of GRBs, as a short pulse of low-frequency radio emission (\\cite{Katz94,Katz99}). ",
        "conclusions": "In this Letter, we have shown that GRBs may be accompanied by very powerful short pulses of low-frequency radio emission. For detection of these radio pulses it may be necessary to perform observations at lower frequencies than are generally used in radio astronomy, which are limited by the problem of transmission through and refraction by the ionosphere.  In particular, observations from space are free of ionospheric refraction and are shielded by the ionosphere from terrestrial interference.  Although even harder to predict, detection from the ground at higher frequencies may also be possible.  Space observations are possible at frequencies down to that at which the interstellar medium becomes optically thick to free-free absorption.  This frequency is \\begin{equation} \\nu_{abs} = 1.0 \\times 10^6 \\left\\langle{n_{0.03}^2 \\over T_3^{3/2}} \\right\\rangle^{1/2} \\vert\\csc b^{II}\\vert^{1/2}\\ {\\rm Hz}, \\label{ISMabs} \\end{equation} where $n_{0.03} \\equiv n_e/(0.03$ cm$^{-3}$) and $T_3 \\equiv T/10^{3\\,\\circ}$K (\\cite{Spitzer62}), $n_e$ and $T$ are the interstellar electron density and temperature, respectively, and $b^{II}$ is the Galactic latitude.  The expression $\\langle\\rangle$ may be $O(1)$, but could be substantially larger if the electrons are strongly clumped or cold. On the other hand, although much of the interstellar volume is filled with very hot ($10^{6\\,\\circ}$K) and radio-transparent gas, this probably contains very little of the electron column density.  Intergalactic absorption poses a somewhat less restrictive condition if the medium is hot ($10^{6\\,\\circ}$K), as generally assumed.  At these low frequencies, and even at tens of MHz, interstellar scintillation (\\cite{Goodman97}) will be very large.  Observations of coherent radio emission from GRB would not only illuminate the physical conditions in their radiating regions, but would determine (through the dispersion measure) the mean intergalactic plasma density and (through the scintillation) its spatial structure.  The flux density implied by Eq. (\\ref{Rflux}) appears impressively large, but it applies only to the brief period when the dispersed signal is sweeping through the bandwidth $\\Delta\\nu$ of observation, so that it is unclear if it is, in fact, excluded by the very limited data (\\cite{Cortiglioni81}) available.  Further, the extrapolation of the radiated spectrum to $\\nu \\gg \\nu_{max}$ is very uncertain.  Finally, we have also not considered the (difficult to estimate) temporal broadening of this brief transient signal by intergalactic scintillation, which will both reduce its amplitude and broaden its time-dependence.  Searches for radio pulses started about 50 years ago, prior to the discovery of GRBs. During wide beam studies of ionospheric scintillations, simultaneous bursts at 45 MHz of $10-20$ s duration were reported by Smith (1950) at sites 160 km apart. These events were detected at night, approximately once a week. The origin of these bursts was never determined. The results of Smith (1950) were not confirmed by subsequent observations at frequencies of 150 MHz or higher.  Modern observations (\\cite{Dessenne96} at 151 MHz; \\cite{Balsano98} at 74 MHz; \\cite{Benz98} broadband; see \\cite{Frail98} for a review) set upper bounds to the brightnesses of some GRB at comparatively high frequencies. These bounds are not stringent, and do not exclude extrapolations of the lower frequency fluxes suggested here.  There are few data at lower frequencies."
    },
    "0002/astro-ph0002093_arXiv.txt": {
        "abstract": "The history of hydrodynamic numerical simulations for accretion disks in close binary systems is reviewed, in which emphasis is placed, in particular, on the facts that spiral shock waves were numerically found in 1986 by researchers including one of the present authors and that spiral structure was discovered in IP Pegasi in 1997 by Steeghs et al.  The results of our two and three-dimensional numerical simulations in recent years are then summarized, with comparison being made with observations. ",
        "introduction": "The term ``accretion'' used in astrophysics means the infall of gas onto a celestial body due to the gravitational attraction of the body.  On this occasion, the gravitational energy of the gas is released and, eventually, transformed into radiation and emitted to space.  The accretion process may be classified into two types depending on the momentum possessed by the gas, i.e. whether the gas has a large angular momentum with respect to the accreting object or not.  In the former case, a thin gaseous disk, ``accretion disk'', is formed around the accreting object.  In such a disk, the gas gradually loses its angular momentum by some mechanism and accretes onto the central object.  On the other hand, a gas with less angular momentum rapidly accretes onto the object directly and this is called ``wind accretion'', a process characterized by, when the wind is supersonic, the formation of ``bow shock'' in the forepart of the accreting object.  The present paper deals with accretion disks.  Accretion disks are further classified into those in which the accreting body has a mass comparable to that of typical star and those in which the accreting body is a giant black hole expected to be present at the central core of galaxy.  Our subject in this paper relates to the former.  With the mass-accreting star being what is known as a compact star, which has a very small size compared to its mass, such as black hole, neutron star or white dwarf, a large amount of energy is released on accretion.  Such accretion disks are present in close binaries, including cataclysmic variables, novae and X-ray stars.  In accretion disks, for gas to accrete onto the central body, the gas must lose, by some mechanism or another, its angular momentum.  The standard model proposed by Shakura \\& Sunyaev (1973) states that the angular momentum is transported from the inner parts of the disk to the outer parts, due to some kind of viscosity (see also Shakura 1972a, b, Pringle \\& Rees 1972). The gas of the inner parts, having lost angular momentum, accretes onto the central star, while a fraction of the outer gas, having obtained larger angular momentum, transports it to infinity.  The gas can thus, almost entirely, accrete, while conserving the total angular momentum.  A generally accepted candidate for this viscosity is turbulent viscosity. The standard model is often called the ``$\\alpha$-disk model'', since the magnitude of the turbulent viscosity has been characrerized by a phenomenological parameter, $\\alpha$.  For accretion disks, the Reynolds number is as large as \\(10^{11-14} \\).  The usual hydrodynamical common sense expects a flow with such a large Reynolds number to be turbulent.  With a rotational fluid, however, the situation is not so simple.  Keplerian disks, with angular momentum increasing outwards, amply satisfy the Rayleigh criterion of stability.  That is, where the gas rotates with nearly Keplerian motion, there is no convincing evidence so far for the accretion disk to become unstable.  In contrast, there is available theoretical and numerical evidence of the stability of Keplerian disks (Balbus, Hawley \\& Stone, 1996). Balbus \\& Hawley (1991) concludes that a small magnetic field present in the disk will be amplified during differential rotation and, eventually, generate magnetic turbulence.  This has become one of the most popular models in recent years.  Apart from the standard model, Sawada, Matsuda \\& Hachisu (1986a, b, 1987) conducted two-dimensional numerical simulations of accretion disk, to discover the presence of spiral shock waves in the disk.  Spruit (1987) found self-similar solutions having spiral shock waves.  They proposed a model in which the spiral shock waves absorb angular momentum from the gas. This is, eventually, transported to the orbital angular momentum of the binary system, due to the axial asymmetry of the density distribution in the disk of the gas and to the torque caused by the tidal force of the companion star. We call this the ``spiral shock model'', and we describe it in detail in the present paper. Many textbooks and reviews of accretion disks are available (for example, Pringle 1981, Frank, King \\& Raine 1992, Spruit 1995, Hartmann 1998, Kato, Fukue \\& Mineshige 1998). ",
        "conclusions": "\\begin{enumerate} \\item Spiral shock waves on accretion discs in close binary systems were found numericaly by Sawada et al. (1986a, b, 1987) in their two-dimensional hydrodynamic finite difference simulation. \\item Spiral shock waves are generated by an oval deformation of the gravitational potential and the Lindblad resonances associating with it. Spiral shocks are seen in galactic disks and the primordial solar nebula as well as accretion disks. \\item Three-dimensional simulations also exhibit the presence of spiral shocks despite of some counterarguments. \\item The stream from L1 point penetrates into the accretion disk rather than being blocked by it. Therefore, so-called hot spot is not formed in our present case,  which fact agrees with the results by Bisikalo et al. (1997a, b, 1998a, b, c). \\item Spiral structure was found in the accretion disk of a dwarf nova IP Pegasi by Steeghs et al. (1997) using Doppler tomography technique. \\item Theoretical Doppler maps are made based on two-dimensional simulations, and they agree well with observed ones despite of high temperature of gas in the simulation. \\end{enumerate}"
    },
    "0002/astro-ph0002436_arXiv.txt": {
        "abstract": "Many of the faint submm sources uncovered by deep SCUBA surveys still remain unidentified at most, if not all other wavelengths. The most successful method so far of obtaining accurate positions and hence allowing more secure identifications has been to use centimetre wavelength imaging with the VLA. This has led to a tempting approach for obtaining redshift estimates for this population (Carilli \\& Yun 1999,2000), which relies on the tight FIR--radio relationship and takes advantage of the steep spectral slope in the submm. In this paper we use the submm data from the SCUBA Local Universe Galaxy Survey (SLUGS) to estimate the usefulness of, and the uncertainties in, the radio--submm redshift estimator. If the submm--radio spectral index were correlated with either dust temperature or 850$\\mu$m luminosity, this method could produce biased redshift estimates for 850$\\mu$m selected galaxies. We find, however, that within SLUGS, these correlations are not significant. The ratio of 850$\\mu$m/1.4 GHz flux was found to decrease with increasing radio and FIR luminosity, and we propose that this is due to a component of the dust not associated with recent star formation, but which is instead heated to 15--20 K by the general interstellar radiation field. ",
        "introduction": "Ever since the discovery of Ultra-luminous Infrared Galaxies (ULIRGs) (Joseph et al. 1985; Sanders et al. 1988) which emit up to 99 per cent of their bolometric luminosity in the far-IR, there have been suspicions that our optical/UV view of the early universe might be biased by large amounts of dust obscuration.  These were supported by the discovery of the Cosmic infrared Background (CIB) (Puget et al. 1996; Fixsen et al. 1998), which contains up to twice the luminosity of the integrated optical/UV background. It thus seems likely that much of the star formation activity in the early universe takes place in obscured environments. Many deep surveys have been conducted with the SCUBA bolometer array on the James Clerk Maxwell Telescope (JCMT)\\footnote{The JCMT is operated by the Joint Astronomy Center on behalf of the UK Particle Physics and Astronomy Research Council, the Netherlands Organization for scientific Research and the Canadian National Research Council.} in order to uncover the sources responsible for this background. Significant numbers of sources, far in excess of no evolution models have been found at both submm (Smail, Ivison \\& Blain 1997; Barger et al. 1998; Hughes et al. 1998; Eales et al. 1999) and far-IR (Puget et al. 1999; Kawara et al. 1998) wavelengths.  The next step after finding these sources is identifying them and determining their redshifts, thus allowing the true star formation history of the universe to be traced. This has proved to be a difficult task, due to the large SCUBA beam size at 850$\\mu$m (15 arcsec) and the several arcsec positional uncertainty inherent in all of the SCUBA maps (Ivison et al. 1998,2000; Downes et al. 2000; Barger, Cowie \\& Richards 2000; Eales et al. 2000). The SCUBA sources can sometimes be identified with several plausible optical candidates, or in other cases they remain optical blank fields, even to the magnitudes probed by HST and Keck. A possible solution to the problem of poor positional accuracy lies in the tight correlation between thermal FIR emission and synchrotron radio emission seen in the local universe (Helou et al. 1985; Condon 1992). The mechanism for this relationship is believed to be the massive stars from recent star-formation, which both heat the dust to produce the FIR flux and, when they explode as supernovae, provide relativistic electrons which constitute the synchrotron radio emission. This means that radio observations with the VLA are also sensitive probes of star formation, albeit not as sensitive as the submm at higher redshifts given the unfavourable K-correction at radio wavelengths. For example, Barger, Cowie \\& Richards (2000) have recently targeted 1.4 GHz sources in the deep VLA images (Richards 1999) of the Hubble Flanking Fields (HFF) with SCUBA and found that the majority of bright ($> 6$ mJy) submm sources have radio counterparts. By observing the submm sources from the deep surveys at radio frequencies (particularly 1.4 GHz), more accurate positions can be obtained, allowing more secure identifications of the submm sources with optical or near IR counterparts. However, even with the accurate VLA positions, many SCUBA sources still have no detectable optical/IR counterparts, or they are too faint for spectroscopy. This creates a great problem in determining the redshifts for the sources, necessary to trace their evolution.  One possible route for obtaining redshift estimates for these objects has been suggested by Carilli \\& Yun (1999). This is based on the FIR-radio correlation plus the break in the spectral slope at around 3 mm, where the thermal dust emission takes over from the declining synchrotron tail. This puts the radio on one side of the spectral break and the FIR (or submm in this case) on the other, creating a redshift sensitive ratio. Carilli \\& Yun (1999) used the submm-radio spectral index, defined as: \\[ \\alpha^{850}_{1.4} = 0.42 \\times \\log\\left(\\frac{S_{850}}{S_{1.4}}\\right) \\] where $S_{850}$ and $S_{1.4}$ are the fluxes at 850$\\mu$m and 1.4 GHz.  The exact dependence of $\\alpha^{850}_{1.4}$ on redshift is sensitive to a few parameters: \\begin{enumerate} \\item{The slope of the thermal Rayleigh-Jeans tail in the FIR and submm. This depends on frequency as $\\nu^{2+\\beta}$ where the dust emissivity index ($\\beta$) is thought to lie between 1 and 2 (Hildebrand 1983). It is the steepness of this slope which potentially makes the indicator sensitive to redshift} \\item{The slope of the radio synchrotron emission. For star forming regions this is assumed to be $-0.7\\,\\Rightarrow \\,-0.8$ (Condon 1992). At higher frequencies this will flatten due to radio free-free emission but this is not simple to model. Free-free absorption at lower frequencies can also lead to a flattening of the radio spectra, the likely effect of this on $\\alpha^{850}_{1.4}$ is discussed in Carilli \\& Yun (2000) and they conclude that it is unlikely to be a dominant cause of scatter in the relationship. In any case the slope at radio frequencies is much shallower than in the submm and any uncertainties in it have less effect.} \\item{The temperature of the dust. This determines the redshift at which the thermal spectrum turns over. When this happens the ratio only depends weakly on redshift, and so becomes useless as a redshift indicator. Accurate dust temperatures are notoriously difficult to estimate and require measurements at many different wavelengths throughout the submm and FIR.} \\item{Contribution by AGN. Galaxies containing a radio-loud AGN will have higher radio emission than a purely star-forming galaxy and so lower values of $\\alpha^{850}_{1.4}$. This could lead to ambiguity between high redshift objects with AGN and lower redshift starbursts.} \\item{The FIR-radio correlation and its variability with redshift. In using this technique to estimate redshifts, an inherent assumption is made that the FIR-radio correlation does not vary with redshift. Possible causes of such a change could be a variation in magnetic field strengths in galaxies in the past, or changes in the dust mass opacity coefficient which relates the FIR/submm emission to the mass of dust. The latter would depend on the dust composition and grain sizes so there is scope for it to have been different at the epoch of galaxy formation, when metallicities would have been lower.} \\item{Linked to (v) is the possibility that the FIR-radio, or more importantly, the submm-radio relationship itself depends on some other galaxy property such as dust temperature, luminosity etc.} \\end{enumerate} The possible impact of points i--vi above will be discussed later.  Several authors have already used this technique to determine redshift estimates or limits for sources discovered in deep submm surveys with SCUBA (Hughes et al. 1998; Lilly et al. 1999; Smail et al. 1999; Barger et al. 2000; Eales et al. 2000). This paper uses data from the SCUBA Local Universe Galaxy Survey (SLUGS) to examine the usefulness of the submm-radio ratio as a redshift indicator, to try to normalise it with respect to the local universe and to estimate the uncertainties in it.  We will use $H_0=75$ km s$^{-1}$ Mpc$^{-1}$ and $q_0 = 0.5$ throughout. ",
        "conclusions": "We have used the submm data from a large, complete sample of local {\\em IRAS\\/} galaxies, along with complementary radio data, to define a median redshift estimator using the change in spectral slope between the submm and the radio at $\\sim 3$mm. The scatter in the data has been used to provide $1 \\sigma$ uncertainties on the relationship. The limited number of submm sources from the deep surveys with both radio fluxes and spectroscopic redshifts generally agree with the estimator (within the uncertainties), except for one object which is known to be an AGN (Fig~\\ref{car2F}). {\\em The estimator is useful in a statistical sense rather than for predicting the redshifts of individual objects\\/} as there are many uncertainties, especially the possible contribution to the radio flux by AGN which would lower the data relative to the models thus leading to an under-estimated redshift. Our redshift indicator differs from the recent one of Carilli \\& Yun (2000) in terms of shape and normalisation. The difference is significant, particularly at redshifts $>2$ (Fig~\\ref{carF}). The discrepancy can be attributed primarily to the larger sample size used in this work and also, in part, to the under-estimation of submm fluxes used in CY 2000. An erratum to CY 2000 has since been published which produces a revised redshift estimator using the corrected Lisenfeld et al. (2000) submm fluxes. This removes $\\sim 30$ per cent of the discrepancy between the estimators, and we postulate that the remaining difference is a result of sample selection (the Lisenfeld sample is not a complete sample and contains a higher proportion of radio bright objects compared to SLUGS), and the smaller sample size (17 objects compared to 104 in SLUGS).  The submm-radio spectral index decreases as radio and FIR luminosity increase, but shows no strong trend with 850$\\mu$m luminosity (Fig.~\\ref{alpha_lumF}). The correlations are believed to be due to the sensitivity of the non-thermal radio and thermal FIR emission to recent star formation, while a significant fraction of the 850$\\mu$m flux may be due to colder dust, and not linked to the star formation rate. The absence of a correlation between $\\alpha^{850}_{1.4}$ and 850$\\mu$m luminosity means that any differences in $L_{850}$ between SLUGS and the SCUBA-selected galaxies will not lead to biased redshift estimates. Importantly, we also find no significant evidence for a systematic variation of $\\alpha^{850}_{1.4}$ with dust temperature (Fig.~\\ref{alpha_tempF}) implying that the redshift-temperature degeneracy hypothesised by Blain (1999) does not play a dominant part in the scatter of the $\\alpha^{850}_{1.4}$ -- redshift relation.  We have investigated the impact of neglecting a possible colder dust component in the SEDs of the local galaxies (Fig~\\ref{greyF}). While sensitive to the parameters assumed, in general the differences are not very significant for most plausible models and the agreement is still within the $1 \\sigma$ uncertainties at lower redshifts ($<2$). If cold components are present but unaccounted for, using the median line in Fig~\\ref{carF} to estimate redshifts will most likely lead to an over-estimate of $z$. \\vspace{0.2cm}  We would like to thank Chris Carilli and Ute Lisenfeld for useful discussions. The support of PPARC is gratefully aknowledged by L. Dunne, S. Eales and D. Clements."
    },
    "0002/astro-ph0002399_arXiv.txt": {
        "abstract": "We report the detection of the Diffuse Interstellar Bands (``$DIBs$'') in the optical spectra of seven far-infrared-selected starburst galaxies. The $\\lambda$6283.9 \\AA\\ and $\\lambda$ 5780.5 \\AA\\ features are detected with equivalent widths of $\\sim$ 0.4 to 1 \\AA\\ and 0.1 to 0.6 \\AA\\ respectively. In the two starbursts with the highest quality spectra (M82 and NGC2146), four other weaker $DIBs$ at $\\lambda$ 5797.0 \\AA, 6010.1 \\AA, 6203.1 \\AA, and 6613.6 \\AA\\ are detected with equivalent widths of $\\sim$ 0.1 \\AA. The region over which the $DIBs$ can be detected ranges from $\\sim$ 1 kpc in the less powerful starbursts, to several kpc in the more powerful ones. The gas producing the $DIBs$ is more kinematically quiescent on-average than the gas producing the strongly-blueshifted $NaI\\lambda\\lambda$5890,5896 absorption in the same starbursts. We show that the $DIBs$ in these intense starbursts are remarkably similar to those in our Galaxy: the relative strengths of the features detected are similar, and the equivalent widths follow the same dependence as Galactic $DIBs$ on $E(B-V)$ and $NaI$ column density. While the ISM in starbursts is heated by a photon and cosmic ray bath that is $\\sim$ 10$^3$ times more intense than in the diffuse ISM of the Milky Way, the gas densities and pressures are also correspondingly larger in starbursts. This ``homology'' may help explain the strikingly similar $DIB$ properties. ",
        "introduction": "The Diffuse Interstellar Bands (``$DIBs$'') have been studied for over 60 years, since Merrill (1934) first established their origin in the interstellar medium. Despite decades of intensive investigation, the identity of the carrier or carriers of the $DIBs$ has not been established (see the comprehensive review by Herbig 1995). The most likely candidates are large carbon-rich molecules (e.g. Sonnentrucker et al 1997), perhaps Polycyclic Aromatic Hydrocarbons (PAH's - Salama et al 1999). The strongest and best-studied $DIBs$ in the optical spectrum are known empirically to trace the $HI$ phase of the ISM, with strengths that correlate well with the line-of-sight color excess $E(B-V)$, the $HI$ column density, and the $NaI$ column density as probed with the $NaI \\lambda\\lambda$5890,5896 (``$NaD$'') doublet (see Herbig 1993).  To date, $DIBs$ have been observed almost exclusively in our own Galaxy, and to a limited extent in the Magellanic Clouds (Morgan 1987). Supernova 1986G in NGC 5128 (Cen A) allowed the detection of $DIBs$ produced within the famous dusty gas disk in this elliptical galaxy (di Serego Alighieri \\& Ponz 1987). Most recently, Gallagher \\& Smith (1999) have reported the possible discovery of the $DIB$ at $\\lambda$6283.9 \\AA\\ in the spectra of two ``super starclusters'' near the nucleus of the prototypical starburst galaxy M 82. This is intriguing, since it suggests that the $DIB$ carriers are present at a normal level even in the ISM of an intense starburst, in which the ambient radiation intensity and gas pressure are orders-of-magnitude higher than in the diffuse gas in the Milky Way disk (e.g. Colbert et al 1999).  We have recently analyzed the properties of the interstellar $NaD$ absorption- line in the spectra of 18 high-luminosity, infrared-selected (dusty) starbursts (Heckman et al. 2000 - hereafter HLSA). In the course of this analysis, we examined the 7 starbursts with the highest quality spectra for the presence of $DIBs$ (section 2). As we report in section 3 below, we have detected one or both of the $\\lambda$6283.9 \\AA\\ and $\\lambda$5780.5 \\AA\\ $DIB$ features (normally the two strongest $DIBs$ in the optical spectral region) in all seven cases. We have also been able to map the spatial distribution of the $DIBs$. These data allow us to directly compare the properties of the $DIBs$ in the ISM of these extreme starbursts to sight-lines in the Galaxy having similar gas column density and reddening (section 4). ",
        "conclusions": "\\subsection{Comparison to Galactic DIBs}  The strengths of the prominent Galactic $DIBs$ correlate well with the column densities of both $HI$ and $NaI$ and with the reddening along the line-of-sight (e.g. Chlewicki et al 1986; Herbig 1993). This implies that the $DIB$ carrier is most plausibly associated with the cool atomic phase of the ISM. We can use the data discussed in HLSA to estimate the values for $N_{NaI}$ and $E(B-V)$ in our sample of seven starbursts, to see if the $DIBs$ in our starburst sample obey the same empirical relations defined by the ISM of the Milky Way.  We follow HLSA and derive estimates for $N_{NaI}$ using the average of the values obtained from the classical ``doublet ratio'' method (Spitzer 1968) and the variant described by Hammann et al (1997). We estimate the line-of-sight reddening to the stellar continuum using the observed colors compared to theoretical models for a starburst stellar population (Leitherer et al 1999; see HLSA for details). We list the results in Table 1.  Our best-measured $DIB$ by-far is the strong $\\lambda$6283.9 feature. The data compiled by Chlewicki et al (1986) and Benvenuti \\& Porceddu (1989) show that the mean ratio of the equivalent width of this feature and the color excess is $<W_{6284}/E(B-V)>$ = 1.2 \\AA\\ for heavily-reddened Galactic sight-lines. For our small starburst sample we find a similar result: $<W_{6284}/E(B-V)>$ = 0.8 \\AA. This is shown in Figure 4 where we have plotted $W_{6284}$ {\\it vs.} $E(B-V)$ for a large sample of Galactic sight-lines using the extensive data compled by Herbig (1993). To compare our starburst data directly to this Galactic data we have converted the values given by Herbig (1993) for the equivalent width of the $\\lambda$5780.5 \\AA\\ $DIB$ into estimated values for $W_{6284}$ assuming that the mean ratio measured by Chlewicki et al (1986) and Benvenuti \\& Porceddu (1989) applies ($W_{6284}/W_{5780}$ = 2.2). In Figure 5 we have likewise plotted $W_{6284}$ {\\it vs.} $N_{NaI}$ for both the Galactic data and our starburst data. The starbursts lie at the high-end of the relationship defined by $DIBs$ in the Milky Way.  \\subsection{Relationship to the $\\lambda$2175 \\AA\\ Dust Feature}  Over the years, there has been considerable speculation as to a possible connection between the $DIBs$ and the strong and broad feature at $\\lambda$2175 \\AA\\ in the Galactic extinction curve (see Benvenuti \\& Porceddu 1989). In this context, the detection of strong $DIBs$ in starburst spectra is noteworthy. As shown by Calzetti et al (1994), the $\\lambda$2175 feature is extremely (undetectably) weak in the UV spectra of starbursts. This implies that the carriers of the $DIBs$ and the $\\lambda$2175 feature must be quite distinct (in agreement with the conclusions of Benvenuti \\& Porceddu (1989) for the Galactic ISM).  \\subsection{Speculations}  On the face of it, the above results may seem surprising given the extreme differences between the physical conditions in the ISM of intense starbursts and our own Galactic disk. The strong starbursts in our sample have bolometric surface brightnesses of $\\Sigma_{bol} \\sim 10^{10}$ to $10^{11}$ L$_{\\odot}$ kpc$^{-2}$ (e.g. Meurer et al. 1997), typical star-formation rates per unit area of $\\Sigma_{SFR} \\sim$ 10 M$_{\\odot}$ year$^{-1}$ kpc$^{-2}$, and surface mass densities in gas and stars of $\\Sigma_{gas} \\sim \\Sigma_{stars} \\sim$ 10$^9$ M$_{\\odot}$ kpc$^{-2}$ (e.g. Kennicutt 1998). These are roughly 10$^3$ ($\\Sigma_{SFR}$), 10$^2$ ($\\Sigma_{gas}$) and 10$^1$ ($\\Sigma_{stars}$) times larger than the corresponding values in the disks of normal galaxies. These values for $\\Sigma_{bol}$ correspond to a radiant energy density inside the star-forming region that is roughly 10$^3$ times the value in the ISM of the Milky Way (and see Colbert et al 1999 for direct measurements of this quantity). The rate of mechanical energy deposition (supernova heating) per unit volume in these starbursts is of-order 10$^3$ times higher than in the ISM of our Galaxy (e.g. Heckman, Armus, \\& Miley 1990), as is the cosmic ray heating rate (Suchkov, Allen, \\& Heckman 1993). Finally, simple considerations of hydrostatic equilibrium imply correspondingly high pressures in the ISM: $P \\sim G \\Sigma_g \\Sigma_{tot} \\sim$ few $\\times$ 10$^{-9}$ dyne cm$^{-2}$ (P/k $\\sim$ few $\\times$ 10$^7$ K cm$^{-3}$, or several thousand times the value in the local ISM in the Milky Way). These high pressures have been confirmed observationally (e.g. Heckman, Armus, \\& Miley 1990; Colbert et al 1999).  The interesting result of the above is that despite the extreme conditions prevailing inside these starbursts, the dimensionless ratio of the ISM pressure to the energy density in UV photons (or cosmic rays) is quite similar in starbursts and the disk of the Milky Way. This would in turn imply that (for a given ISM temperature) the ratio of the number densities of the gas particles and UV photons (or cosmic rays) would also be similar to their values in the local ISM. Wang, Heckman, \\& Lehnert (1998) have discussed the evidence that this analysis is correct for the diffuse ionized medium in starbursts and the disks of normal late-type galaxies.  This ``homologous'' behavior of the ISM in regions spanning over three orders-of-magnitude in heating and cooling rates per particle may help to explain why the ratio of the column density of $DIB$ carriers to that of both $Na$ atoms (Figure 5) and dust grains (Figure 4) appears so similar in extreme starbursts and the ISM of our own Galaxy. In the absence of a well-understood origin for the $DIBs$, further speculation seems premature."
    },
    "0002/astro-ph0002350_arXiv.txt": {
        "abstract": "We present evidence that there is a significant correlation between the fraction of a galaxy's mass that lies in its central black hole and the age of the galaxy's stellar population.  Since the absorption line indices that are used to estimate the age are luminosity weighted, they essentially measure the time since the last significant episode of star formation in the galaxy.  The existence of this correlation is consistent with several theories of galaxy formation, including the currently-favoured hierarchical picture of galaxy evolution, which predicts just such a relation between black hole mass and the time since the last burst of merger-induced star formation. It is not consistent with models in which the massive black hole is primordial, and hence uncoupled from the stellar properties of the galaxy. ",
        "introduction": "The existence of active galactic nuclei has long been taken as evidence for the existence of massive black holes in the centres of some galaxies (Lynden-Bell 1969).  However, it is only relatively recently that high spatial resolution studies of the kinematics of galactic nuclei have revealed that essentially all galaxies harbour large central masses [see Ho (1999) for a review of the evidence]. The existence of these observations also means that there are now enough data to study the demographics of massive black holes, in order to seek clues to their origins.  The first significant discovery in this regard is that there is a correlation between the mass of the black hole, $M_{\\rm BH}$, and the mass of the host galaxy's spheroidal component, \\footnote{The term ``spheroidal component'' refers to the whole system in the case of elliptical galaxies, but just the bulge in systems with significant disk components.} $M_{\\rm sph}$.  Although there is a variety of possible biasses in measuring this correlation, it seems broadly to be the case that there is a linear relationship, such that $M_{\\rm BH} \\sim 0.005 M_{\\rm sph}$ (Magorrian et al.\\ 1998).  Although this correlation is reasonably strong, there is still considerable scatter in the relation, such that there is more than a factor of ten variation in the inferred value of $M_{\\rm BH}$ for galaxies of given spheroid mass (Magorrian et al.\\ 1998).  Some of this scatter can probably be attributed to the uncertainties in calculating black hole masses from relatively poor kinematic data and simplified dynamical models (van der Marel 1997).  However, there are also astrophysical reasons why one might expect significant dispersion in this relation.  For example, consider the simplest possible scenario in which galaxies form and evolve in near isolation.  If the central black holes in these galaxies accrete mass fairly steadily from their hosts, then the mass of a black hole simply reflects the age of its host.  Under the currently-favoured hierarchical paradigm for galaxy formation, in which larger galaxies are formed from the merging of smaller galaxies (White \\& Rees 1978), the simple linear correlation between galaxy mass and black hole mass is readily explained.  Each time two galaxies merge to form a larger system, their black holes rapidly spiral to the centre of the new galaxy due to dynamical friction.  The black holes then merge, creating a proportionately-larger black hole.  However, a galaxy formed by this process of repeated mergers cannot be characterized by a single age, so the above explanation for the scatter in black hole masses must be modified somewhat.  One measure of such a galaxy's age is the time since it last underwent a major merger, and Kauffmann \\& Haehnelt (2000) have shown that this timescale is a key factor in explaining the scatter in black hole masses.  If the last merger happened long ago, then it will have occurred between relatively unevolved galaxies in which there would have been a large amount of cold gas.  If the black hole accretes some fixed fraction of this gas, then galaxies in which the last merger occurred longer ago will contain more massive black holes.  This picture, in which black holes acquire much of their mass through accretion of material from their host galaxies, seems quite credible. However, it is not the only possible scenario.  Stiavelli (1998) has argued that galaxies with essentially identical properties could be formed around pre-existing massive black holes.  In this case, the spread in black hole masses would simply reflect the stochastic nature of whatever physics was responsible for the formation of the primordial black holes.  In this Letter, we investigate whether we can attribute the observed scatter in black hole masses to an astrophysical cause, and hence whether we can discriminate between the above scenarios. Specifically, we investigate whether the masses of black holes correlate with the ages of their host galaxies as determined by stellar absorption line diagnostics. ",
        "conclusions": "Although there does appear to be a significant correlation between measured black hole mass and galaxy age estimate, it is not necessarily astrophysical in origin.  We must first consider the possibility that it arises from some systematic error in the analysis. However, the kinematic data from which the black hole masses were inferred are completely independent from the line index data that provide the age estimates.  Since the line index data were not selected with this project in mind, and the black hole mass estimates played no role in the choice of sample, the selection process cannot have induced the correlation that is seen in Fig.~1.  Further, the independent nature of the data sets used to measure the two ordinates means that there can be nothing in this analysis that might preferentially over-estimate the black hole masses in old galaxies, or underestimate the masses in young systems.  It should also be borne in mind that the absolute calibrations of the black hole masses and galaxy ages are significantly uncertain.  In the case of the absorption line indices, for example, the age estimates are derived from spectral synthesis modelling, which remains a somewhat uncertain process, so the absolute values of the ages of two galaxies may be quite ill-determined.  However, the fact that one is older than the other can be determined relatively reliably by this modelling process, so the approximate ordering of galaxy ages can be determined quite robustly.  Since the Spearman rank test described above depends only on this ordering, the statistical significance of the correlation is not dependent on the details of the adopted calibration.  It would thus appear that there is an underlying astrophysical correlation between the fraction of a galaxy's mass in its central black hole and the age of its most recently formed stellar component. Hence, in addition to the established correlation between black hole mass, $M_{\\rm BH}$ and galaxy mass, $M_{\\rm sph}$, there seems to be a ``second parameter'' correlation with the age of the youngest stellar component.  At any given value of $M_{\\rm sph}$, different age galaxies will have different values of $M_{\\rm BH}$, so this secondary correlation must go some way toward explaining the scatter in the primary relation.  We have sought to quantify the contribution of this second parameter to the scatter in the relation between $M_{\\rm sph}$ and $M_{\\rm BH}$ by calculating \\begin{equation} \\log(M_{\\rm BH}/M_{\\rm sph})^* = \\log(M_{\\rm BH}/M_{\\rm sph}) - \\log(t/10\\,{\\rm Gyr}). \\end{equation} This process corrects the mass ratio for the effects of age by subtracting the simplest possible linear fit to the correlation in Fig.~1.  As one would expect, this correction reduces the scatter in the relation: for the data in this sample, the dispersion in $\\log(M_{\\rm BH}/M_{\\rm sph})$ is 0.42 dex while that in $\\log(M_{\\rm BH}/M_{\\rm sph})^*$ is only 0.31 dex.  Clearly, even the corrected mass ratio still contains considerable scatter.  However, given the large uncertainties in the individual black hole mass and galaxy age determinations, it would be very surprising if the dispersion were reduced below a factor of two ($\\sim 0.3$ dex).  The simplest explanation for the existence of the second parameter correlation is that a single physical process couples the growth of the central black hole to the triggering of star formation in a galaxy.  As outlined in the Introduction, the hierachical picture of galaxy and black hole evolution described by Kauffmann \\& Haehnelt (2000) suggests that galaxy mergers lie behind both processes.  Where the last major merger occurred long ago, it will have taken place in a gas-rich environment that will provide ample fuel to augment the mass of the black hole.  Since the last major episode of star formation will also be triggered in the merger, such galaxies will contain old stellar populations and massive black holes.  Conversely, galaxies formed in more recent mergers will contain under-massive black holes and younger stellar populations.  Although the correlation between black hole mass and galaxy age is predicted by the hierarchical merging models, it should be borne in mind that such a correlation is a fairly generic prediction of any model in which the black hole mass grows over time.  Even if galaxies form monolithically, those that form first -- and hence contain the oldest stellar populations -- will have had time to grow the largest black holes.  The models that do not fit easily with this correlation are those in which the black holes and stellar components form at entirely different times -- it would be hard to explain the observed correlation if, for example, the central black holes were entirely primordial.  The study of black hole demographics is maturing rapidly, and, as we hope we have shown, it is already possible to detect phenomena beyond the basic relation between $M_{\\rm BH}$ and $M_{\\rm sph}$.  In the near future, larger sets of both the kinematic and line-strength data will become available, and more sophisticated modelling techniques will be developed to refine the estimates for black hole masses and galaxy ages.  With these tools, it will become possible to address subtler questions, such as whether a galaxy's environment plays a significant role in its black hole growth rate.  Such analyses will provide key tests for theories of black hole formation within the broader context of galaxy evolution."
    },
    "0002/astro-ph0002166_arXiv.txt": {
        "abstract": "We present a high resolution \\Chandra\\ X-ray observation of \\psr, the Crab-like 50 msec pulsar in the Large Magellanic Cloud. We use phase-resolved imaging to decompose the extended X-ray emission, as expected of a synchrotron nebula, from the point-like emission of the pulsar.  The image of the pulsed X-ray emission shows a well-defined point-spread function of the observation, while the resolved nebula has a morphology and size remarkably similar to the Crab nebula, including evidence for a jet-like feature from \\psr.  The patchy outer shell, which most likely represents the expanding blast-wave of the supernova, is reminiscent of that seen in radio. Based on morphology, size, and energetics there can be little doubt that \\snr\\ is an analogous system to the Crab but located in our neighboring galaxy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002485_arXiv.txt": {
        "abstract": "High-Velocity Clouds (HVCs) have radial velocities that cannot be explained by the global Galactic rotation; their distances remain mostly unknown, and their true nature and origin are still a mystery. Some of them could be of galactic origin, or belong to tidal streams drawn by the Milky-Way/ Magellanic Clouds interaction, or could even be intergalactic clouds infalling onto the Local Group. In the latter hypothesis, they play a major role in the hierarchical formation scenario of the Milky-Way and are connected to the Lyman-limit absorption systems. In any case, the determination of their physical state (density, temperature, internal structure, abundances, excitation) will help to discriminate between current theories on their origin and nature. A recent UV measurement (Richter et al 1999) has discovered for the first time in a HVC the molecular phase that was previously searched for, without success, through CO emission. Previous non detections could be due either to metallicity problems, or insufficient excitation (because of low density). Low-excitation molecular gas may, however, be detectable though absorption. Here we report on a sensitive search for \\hcoplus(1-0) absorption lines in front of 27 quasars, already known to be strong millimetric continuum sources. Except for one tentative case, no detection was obtained in most HVCs, although \\hcoplus(1-0) was clearly detected towards galactic low-velocity clouds. We discuss the implications of this result. ",
        "introduction": "Since their discovery by Muller, Oort, \\& Raimond (1963), HI clouds moving with velocities that can not be explained by differential Galactic rotation (often exceeding it by 100 \\kms) have been the target of numerous surveys.  Many of these high-velocity clouds (HVCs) are located at intermediate and high Galactic latitudes (Giovanelli, Verschuur, \\& Cram 1973; Mathewson, Schwarz, \\& Murray 1977; Wakker \\& van Woerden 1991) and do not appear to have any connection with the gas in the Galactic disk. HVCs are often found in large HI complexes with angular sizes 10--90$^\\circ$. They cover from 10 to 37\\% of the sky, depending on the sensitivity of the studies (Murphy et al 1995).  The origin of HVCs is still unclear, mainly because the distances to the individual complexes are in most cases unknown.  They could be cold gas corresponding to the return flow in a Galactic fountain (e.g. Houck \\& Bregman 1990), or gas left over from the formation of the Galaxy. Some HVCs are probably belonging to the tidal gas streams torn from the Magellanic Clouds by the Milky-Way (Mathewson et al 1974, Putman \\& Gibson, 1999).  Several authors have already explored the possibility that HVCs are infalling primordial gas and have associated them with the Local Group (see Wakker \\& van Woerden 1997 for a thorough review). Recently Blitz et al. (1999) re-examined this hypothesis, and simulating the dynamical evolution of the Local Group galaxies, used the up-to-date HI maps of HVCs to show that the HVCs are consistent with a dynamical model of infall of the ISM onto the Local group. As such, they would represent the building blocks of our galaxies in the Local group and provide fuel for star formation in the disk of the Milky Way.  In their model, HVCs contain altogether 10$^{11}$ M$_\\odot$ of neutral gas.  From their stability analysis, they conclude that there is roughly 10 times more dark matter than luminous gas within each HVC, and that these could correspond to the mini-halos which are able to accumulate baryons, and can gather into filaments (e.g. Bond et al 1988, Babul \\& Rees 1992, Kepner et al 1997). HVCs would therefore be related to the hierarchical structure of the Universe (see e.g. Katz et al 1996), and to the gas seen in absorption towards quasars (Lyman-$\\alpha$ forest and Lyman-limit lines). However, Giovanelli (1981) has pointed out that the velocity distribution of HVCs does not match that of the Local Group, but does favor an association with the Magellanic Stream, the most obvious tidal feature of the interaction between the Galaxy and the Magellanic Clouds. The discrepancy with the results from Blitz et al. (1999) comes from the fact that the latter authors have not considered all observed HVCs, but only a selection of them.  Maps of the brightest HVC complexes have revealed the existence of unresolved structure at 10 arcmin resolution, which was further resolved into high-density cloud cores at 1 arcmin resolution (Giovanelli \\& Haynes 1977; Wakker \\& Schwarz 1991; Wakker \\& van Woerden 1991).  More generally, HVCs follow the fractal structure observed in the whole interstellar medium (Vogelaar \\& Wakker, 1994). The HI column densities in these cores are estimated to be several times 10$^{20}$ cm$^{-2}$ and their temperatures are generally between 30 and 300 K.  The central densities of individual clouds can reach $>$ 80 cm$^{-3}$ D$_{kpc}^{-1}$, where D$_{kpc}$ is their distance in kpc. Depending on the actual distance, which still remains poorly determined, those conditions make the HVC cores possible sites of star formation. HVCs have in fact been considered good candidates for the source of young Population I stars at large distances from the Galactic plane (see e.g. Sasselov 1993).  Attempts to measure the spin temperature of atomic hydrogen through 21cm absorption in front of background continuum sources have often only resulted in upper limits (Colgan et al. 1990, Mebold et al. 1991). A few detections have been reported (Payne et al. 1980, Wakker et al. 1991, Akeson \\& Blitz, 1999), with inferred spin temperature as low as 36\\,K, but in general most HVCs must have spin temperatures larger than 200\\,K or be very clumpy.  Most of our current knowledge of the gaseous content of HVCs comes from HI observations. Efforts to search for molecular hydrogen using CO emission lines have been so far unsuccessful (Wakker et al 1997), since the sub-solar metallicity and/or low density of HVCs makes direct CO emission line detection very difficult. However, optical absorption lines have shown that HVCs are not completely devoid of heavy elements (e.g. Robertson et al 1991, Lu et al 1994, Keenan et al 1995, Wakker \\& van Woerden 1997).  The lines detected are from SiII, CII, FeII, or CIV, but the strongest are from MgII (Savage et al. 1993, Bowen \\& Blades 1993, Sembach et al. 1995, 1998). Metallicity studies have been done in HVCs to determine their origin. If HVCs result from the Galactic fountain effect, their metallicity should be solar while if they were associated with the Magellanic or intergalactic Clouds, it could be even less than 0.1 solar. In fact, the determined abundances are around 0.1 solar, but with much uncertainties, because of saturated lines, or dust depletion (Sembach \\& Savage 1996).  This average metallicity is compatible with a Local-Group infall model, since X-ray observations have revealed abundances of 0.1 solar in poor groups (Davis et al 1996), and even 0.3 solar in intra-cluster gas (Renzini 1997).  If HVCs were local analogues of Lyman-limit absorbing clouds, background QSOs could enable us to detect molecules in absorption, a task which is considerably easier to achieve.  Absorption has recently been reconfirmed as a very powerful tool in the millimetric range (Lucas \\& Liszt 1996 hereafter LL96, Combes \\& Wiklind 1996). A mm molecular detection would advance our knowledge of HVCs, their physical conditions, and their possible origin.  A first attempt has been made in this domain by Akeson \\& Blitz (1999) with only negative results, using the BIMA and OVRO interferometers.  Here we report on \\hcoplus(1-0) absorption line observations, made in the southern sky with the single dish 15m ESO-SEST telescope, and in the northern sky with the IRAM interferometer. The choice of \\hcoplus~  is justified because, due to its large dipole moment, it is generally not excited in diffuse media (the critical density for excitation is $\\sim$ 10$^7$ cm$^{-3}$); therefore confusion with emission is not a problem, contrary to the CO(1-0) line which requires an interferometer to resolve out the emission. Furthermore, the \\hcoplus~ absorption survey carried out by Lucas \\& Liszt (1994, 1996) has revealed more and wider absorption lines than in CO, suggesting that it might be a better tracer. In that study, the derived abundance \\hcoplus/H$_2$ was surprisingly large, of the order of 6$\\times$10$^{-9}$, and sometimes even an order of magnitude larger.  The details of the observations are presented in Section 2.  Section 3 summarizes the results, which are then discussed in Section 4.    \\begin{figure*}[ht] \\includegraphics[width=\\hsize]{9500.f1} \\caption{ a) All-sky map of the high-velocity clouds, in HI 21cm line, in Aitoff projection, from Wakker \\& van Woerden (1997).  b) Location in the same projection of the quasars observed in this work.} \\label{map} \\end{figure*} ",
        "conclusions": "\\subsection{ The diverse nature of HVCs}  It is still difficult to tackle the problem of the origin of HVCs, since they may be diverse in nature, with possibly different formation histories. The first idea that they are clouds infalling towards the Milky Way (Oort, 1966) was proposed since only negative velocities clouds were discovered at this epoch. However, since differential galactic rotation makes an important contribution to those velocities and the tangential velocity of the HVCs is also unknown, it is difficult to ascertain their true three dimensional motion. Now that almost an equal number of clouds are detected with positive high velocities, and that is also true in the various reference frames (e.g. Wakker 1991), it is believed there must be other explanations for at least some of them.  Some HVC complexes have been proven to be associated with the Galaxy, in being very close (less than 10 kpc) due to absorption detected in front of Galactic halo stars (e.g. Danly et al. 1993, van Woerden et al. 1999).  This supported a Galactic origin for some of the HVCs and more specifically the Galactic fountain model, where gas is ejected in the halo by the star formation feedback mechanisms.  The main problems with this scenario are the low metallicity observed in most lines of sight (even though it could explain why clouds have nearly solar metallicity, e.g. Richter et al 1999), and the extreme velocities sometimes reached (up to -450\\,\\kms~ in the negative side). Alternatively, HVCs could be transient structures, now becoming bound to the Galaxy: an ensemble or complex of HVCs has been identified as tidal debris from the Magellanic Clouds (Mathewson et al. 1974), and is called the Magellanic Stream. Few of them could come from tidal debris from interactions with other dwarf galaxies but the corresponding dwarf galaxies either have not been identified yet, or they have been already dispersed (e.g. Mirabel \\& Cohen 1979).  A key point is to consider whether HVCs are transient structures or coherent and self-gravitating. With only their column density observed in HI, to be self-gravitating they should be at least at 1\\,Mpc distance, and in average at 10 Mpc, therefore outside the Local Group (Oort, 1966). This distance can be reduced by a factor 10, if the total mass of HVC is taken to be 10 times their HI mass, composed mostly of dark matter (Blitz et al. 1999). Since the specific kinetic energy is in M/R $\\propto$ Distance, then the clouds would have to be only at about 1\\,Mpc, and would consequently be part of the Local Group.  The main difficulty with this model of mini-haloes merging with our Galaxy, is that other high-velocity clouds should have been observed around or in between external galaxies, which is not the case (Giovanelli 1981, Zwaan et al. 1997, Banks et al. 1999).  Moreover, the derived typical dimensions and mass of these systems do not correspond to any observed object or any extrapolation of mass ranges (they have dwarf masses, but are very extended in size).  Another possibility, however, could be that HVCs represent gas left over from the formation of the Galaxy: they form a system out of equilibrium, but bound to the Galaxy, and are now raining down to the Galactic disk. This would explain several observed characteristics, such as the 4-10 kpc distance determined for some complexes, and the low metallicity of many of them (e.g. Wakker et al. 1999).  \\begin{figure}[ht] \\includegraphics[width=\\hsize]{9500.f4}  \\caption{ Spin temperature T$_s$ of the HI in HVCs derived from absorption/emission measurements, as a function of the optical depth $\\tau$ of the absorption. Positive detections are from Payne et al. 1980 (filled stars) and Wakker et al. 1991 (unfilled circles). Upper limits in HI absorption lead to lower limits in T$_s$ and are plotted by vertical bars (Payne et al. 1980, Wakker et al. 1991, Akeson \\& Blitz 1999). The line is that fitted to the detections of low-velocity gas (Payne et al. 1983).  } \\label{ts} \\end{figure}    \\subsection { How many molecular absorptions could be expected? }   Given the low HI column densities reported in Table \\ref{tab2}, it could seem hopeless to detect molecular absorption, with the present sensitivities. But this is true only if the HVC gas is homogeneous. In fact, from previous HI observations, we expect that like the gas in the galactic disk,the HVC gas is also composed  of several components, with clumps of much higher column densities, only detectable at higher spatial resolution. This information can be deduced from existing  HI observations in emission and absorption, towards HVCs. The optical thickness $\\tau$ is related to the column density N(HI) (cm$^{-2}$), the spin temperature T$_s$(K) and the FWHM velocity width $\\Delta V$ (\\kms) by $$ \\tau = \\frac {N(HI)}{1.8\\times10^{20}} \\,  \\left(\\frac{50}{T_s}\\right) \\, \\left(\\frac{20}{\\Delta V}\\right) $$ Since the profile extends at least over 20\\,\\kms~ in emission, and the spin temperature has been determined to be at least 50\\,K, if there was only one component, the derived optical depth would have to be very low, $\\tau < 0.1$. Absorption and emission studies in the Galactic plane have shown though that there is indeed more than one component. If there was only one, both emission and absorption profiles would look similar, which is not what is observed. In fact, there are at least two components; a warm diffuse inter-clump component with cloudlets which are cold, narrow in velocity, and more optically thick (Garwood \\& Dickey 1989). If the surface filling factor of the two components are f$_1$ and f$_2$, for spin temperatures T$_1$ and T$_2$, and optical depths $\\tau_1$ and $\\tau_2$, the observed ratio between the absorption depth $\\frac{T_{abs}}{T_{cont}}$ and the antenna temperature of the emission $T_{em}$ is $$ \\frac{T_{abs}}{T_{cont} T_{em}} = \\frac{f_1 (1-e^{-\\tau_1}) +f_2 (1-e^{-\\tau_2}) } {T_1 f_1 (1-e^{-\\tau_1}) + T_2 f_2 (1-e^{-\\tau_2}) } $$ In this formula, it is obvious that each component is weighted according to its mass if it is optically thin (i.e. weight $\\propto f \\tau$), but considerably less if $\\tau >> 1$. Therefore, if the cold component is optically thick, the derived spin temperature will be overestimated\\footnote{In fact the mass of the cold component would be itself under-estimated, since the absorption often saturates, while the emission, dominated by the warm gas, saturates less}.  Are HVC clouds of the same nature than normal low-velocity galactic clouds, and have they similar  small-scale structure? For normal HI clouds, Payne et al. (1983) have determined (from absorption/emission comparisons) that the weighted average spin temperature is decreasing as the optical depth increases (cf line in  Figure \\ref{ts}). There is only a small scatter in this relation, which means that if the cold and thick component is made of cloudlets, their size must be smaller than that of their background continuum sources, i.e. a fraction of an arcmin, and their number must be large accordingly (N $>$ 100). We do not see the situation where absorption features are deep and rare (statistically the T$_s$-$\\tau$ relation will still hold, but with large scatter). The data favor a model, in which the cloudlets are quite small ($<$ 0.1 pc) and numerous, and they are weighted according to their surface filling factor $f$ both in emission and in absorption for any line of sight. The existence of such small scale structure is also confirmed by VLBI HI absorption (Faison et al. 1998), where sizes down to $\\sim$ 20 AU are detected.  It is also very likely that the structure, apart from these smallest fragments, has no particular scale.  If it is a fractal, statistically, there is a correlation between the optical depth between scales with a limited scatter. This would explain the observed correlation between emission and absorption at different scales.  When the emission/absorption measurements for HVCs are considered (cf Figure \\ref{ts}) the lower limits on derived spin temperature are compatible with the detections on the normal low-velocity gas and their T$_s$-$\\tau$ relation. The spin temperature of HVCs is a weighted mean of the warm and cold components, with the same mixtures as the one observed in normal low-velocity clouds of the galactic plane. The difficulty to find HI in absorption in HVCs corresponds then only to their low average column density, and not to a different physical structure.    As for molecular absorptions, the column densities to which a detection is possible is even larger. The corresponding scales must be smaller, and consequently the corresponding line-widths narrower. This is observed for low-velocity absorptions, where the detected line-widths are as narrow as 0.6\\,\\kms (LL96). The detected optical depths are larger on average, and the probability to underestimate the column density is higher, because of saturation. (One should note though that the apparent optical depth is low because of spatial and velocity dilutions.) The probability of detection of the \\hcoplus(1-0) absorption has been estimated to be 30\\% as large as the 21cm absorption for galactic clouds (LL96).  The presently observed low probability to find an \\hcoplus(1-0) absorption is therefore expected. In addition, there could be a column density threshold for self-shielding against photo-dissociation, that hampers molecular observations.  This threshold has been estimated for CO emission to N(H$_2$) $\\sim$ 4$\\times$10$^{20}$\\,cm$^{-2}$ (or equivalently of the order of 10$^{12}$\\,cm$^{-2}$ for \\hcoplus). If our tentative detection is confirmed, given the low metallicity of the HVCs, this will support the existence of clumps of high column densities in this medium. The fact that molecular absorption is more frequent with respect to emission than atomic absorption is related to the excitation mechanism. Large volumic critical densities are required to excite molecules above the cosmic background, and this is in particularly true for \\hcoplus. This makes absorption techniques the more promising to probe the molecular content of HVCs in the future.   \\subsection { Physical nature and distance of the gas}  Since the HVCs appear to follow in projection the same fractal properties as the normal low-velocity gas (Vogelaar \\& Wakker 1994), it could be interesting to develop an insight in their distance from their size-linewidth relation. It is now well established that clouds in the interstellar medium (either molecular or atomic) are distributed according to a self-similar hierarchical structure, characteristic of a fractal structure (Falgarone et al. 1991; Stanimirovic et al. 1999; Westpfahl et al. 1999).  Such a scaling relation between mass and size, can also lead to a relation between size and line-width, provided that the structures are virialized. In particular, various clouds at all scales obey a power-law relation between sizes $R$ and line-widths or velocity dispersion $\\sigma$: $$ \\sigma \\propto R^q $$ with $q$ between 0.35 and 0.5 (e.g. Larson 1981, Scalo 1985, Solomon et al 1987).  We have plotted this latter relation, together with the sizes and line-widths of the 65 well defined isolated HVCs, catalogued by Braun \\& Burton (1999). This catalog is expected to be free from galactic contamination as well as from blending along the line of sight, because of the isolation criterium imposed for their selection.  To compute their sizes, two distances were assumed, either 20 kpc, or 1 Mpc, and the geometrical mean between major and minor axis was calculated. At a distance of 20\\,kpc, the clouds fall on the relation corresponding to Giant Clouds in the Galaxy. Note that the large scatter is due to the fact that this choice of distance is only one order of magnitude, and the HVCs are certainly not all at the same distance.  This set of clouds has been determined with a spatial resolution of half a degree.  The clouds determined at higher spatial resolution have correspondingly narrower line-widths. In Figure \\ref{bb}, we have also plotted the characteristics of clouds determined with the Westerbork interferometer, at 1 arcmin resolution (Wakker \\& Schwarz 1991).  They also fit the galactic clouds relation on average, if their distance is chosen to be 20 kpc, with large scatter since individual distances are also not known. In fact, HI observations at different spatial resolutions emphasize a particular scale of the hierarchical structure (see e.g the 10' resolution observations by Giovanelli \\& Haynes 1977). This hierarchical structure is similar to what is observed for ``normal'' galactic clouds, in the sense that the same fraction 20-30\\% of the single dish flux is retrieved in the interferometer data (Wakker \\& Schwarz 1991).  The size-linewidth relation has been widely observed up to 100pc in size, the largest size for self-gravitating clouds in the Galaxy, but it might appear questionable to extend it to higher scales, where the gas would not be self-gravitating.  At these scales, the gas is bound into largest self-gravitating structures, including stars or dark matter. However, even at these scales, the gas should trace the gravitational potential of the bound structure it is embedded in, and share the corresponding velocity dispersion; such a relation is observed for instance in the form of the Tully-Fisher relation in galaxies. The main point is that the gas should reveal velocity profiles in emission that should grow wider with the distance, if it belongs to an assumed self-gravitating remote system. The observed profile width is therefore a distance indicator.   \\begin{figure}[ht] \\includegraphics[width=\\hsize]{9500.f5}  \\caption{ Size-linewidth plot for HVCs, observed with two spatial resolutions. The first set is the compact isolated HVCs from the Braun \\& Burton (1999) catalog, with two assumed distances, 20 kpc (filled stars), and 1 Mpc (open circles). The second set is the clouds identified by Wakker \\& Schwarz (1991) with the Westerbork interferometer, with the same two assumed distances, 20 kpc (filled triangles), and 1 Mpc (open triangles). The line is the relation derived for clouds in the Milky Way (Solomon et al. 1987).  } \\label{bb} \\end{figure}   \\subsection {Conclusion}   We have searched for \\hcoplus(1-0) absorption towards 27 high velocity clouds, in front of remote radio-loud quasars. The technique is efficient, since we detect the existing absorption due to low velocity galactic clouds, for our low latitude sources. Only one tentative HVC detection is reported. If confirmed, this indicates the presence of small scale cloudlets, at low excitation and high column densities, prolonging the hierarchical structure already observed in the atomic component at larger scale. When this hierarchical structure is compared to the one observed for low-velocity galactic clouds, a good fit is obtained for the self-similar relation between sizes and line-widths, if the HVCs are on average at 20\\,kpc distance.  Since mm molecular absorptions are expected to be more frequent than emission for these low column density HVCs, this absorption technique appears promising to probe the molecular component of HVCs, already directly detected by UV H$_2$ absorption lines (Richter et al 1999)."
    },
    "0002/astro-ph0002216_arXiv.txt": {
        "abstract": "We present VLBI Space Observatory Programme (VSOP) observations of the gamma-ray blazar 3C~279 at 1.6 and 5~GHz made on 1998 January 9-10 with the HALCA satellite and ground arrays including the Very Long Baseline Array (VLBA). The combination of the VSOP and VLBA-only images at these two frequencies maps the jet structure on scales from 1 to 100~mas.  On small angular scales the structure is dominated by the quasar core and the bright secondary component `C4' located 3 milliarcseconds from the core (at this epoch) at a position angle of $-$115$^\\circ$. On larger angular scales the structure is dominated by a jet extending to the southwest, which at the largest scale seen in these images connects with the smallest scale structure seen in VLA images. We have exploited two of the main strengths of VSOP: the ability to obtain matched-resolution images to ground-based images at higher frequencies and the ability to measure high brightness temperatures.  A spectral index map was made by combining the VSOP 1.6~GHz image with a matched-resolution VLBA-only image at 5~GHz from our VSOP observation on the following day. The spectral index map shows the core to have a highly inverted spectrum, with some areas having a spectral index approaching the limiting value for synchrotron self-absorbed radiation of $\\alpha = +2.5$ (where $S \\propto \\nu^{+\\alpha}$). Gaussian model fits to the VSOP visibilities revealed high brightness temperatures ($>10^{12}$~K) that are difficult to measure with ground-only arrays.  An extensive error analysis was performed on the brightness temperature measurements.  Most components did not have measurable brightness temperature upper limits, but lower limits were measured as high as $5~\\times~10^{12}$~K. This lower limit is significantly above both the nominal inverse Compton and equipartition brightness temperature limits.  The derived Doppler factor, Lorentz factor, and angle to the line-of-sight in the case of the equipartition limit are at the upper end of the range of expected values for EGRET blazars. ",
        "introduction": "\\label{intro} The quasar 3C~279 ($z$=0.536) is one of the most intensively studied quasars for several reasons. It was the first radio source observed to exhibit superluminal motion (Knight et al.\\ 1971; Whitney et al.\\ 1971; Cohen et al.\\ 1971).  It was the first blazar --- and remains one of the brightest --- detected in high-energy $\\gamma$-rays by the EGRET instrument on the {\\em Compton Gamma Ray Observatory} (Hartman et al.\\ 1992).  The EGRET detection prompted several large multiwavelength studies of this source.  Results of these studies are presented by Maraschi et al.\\ (1994), Grandi et al.\\ (1996), Hartman et al.\\ (1996), and Wehrle et al.\\ (1998). The radio flux density of 3C~279 has been monitored by the Michigan group since 1965 at frequencies of 4.8, 8.0, and 14.5 GHz (Aller et al.\\ 1985). The observations presented in this paper occurred at the beginning of a total flux density flare that would later reach the highest flux densities yet recorded in this program for 3C~279. \\footnote{http://www.astro.lsa.umich.edu/obs/radiotel/gif/1253-055.gif}  Following the discovery of superluminal motion, the parsec scale structure of 3C~279 has been monitored using the VLBI technique. Cotton et al.\\ (1979) measured a speed of 15~$c$ for the original superluminal jet component. (Throughout the paper we assume $H_{0}$=70 km s$^{-1}$ Mpc$^{-1}$ and $q_{0}$=0.1, and component speeds measured by others have been expressed in these terms.) Unwin et al.\\ (1989) and Carrara et al.\\ (1993) describe VLBI monitoring of 3C~279 at 5, 11, and 22~GHz throughout the 1980s. These authors observed the motions of several new superluminal components, and found that the speeds of these components were only one-quarter to one-third (3-5~$c$) of that measured for the original superluminal component during the 1970s. VLBI monitoring of 3C~279 has been undertaken at 22 and 43~GHz during the 1990s.  Initial results of this high-frequency monitoring are reported by Wehrle, Unwin, \\& Zook (1994) and Wehrle et al.\\ (1996), and final results will be reported by Wehrle et al.\\ (2000). High-resolution VLBI polarimetric images of 3C~279 have been made by Lepp\\\"{a}nen, Zensus, \\& Diamond (1995), Lister, Marscher, \\& Gear (1998) and Wardle et al.\\ (1998). The detection of circularly polarized radio emission by Wardle et al.\\ provides some of the first direct evidence that electron-positron pairs are an important component of the jet plasma. One notable feature of the VLBI observations of this source has been the differing speeds and position angles of the VLBI components.  Carrara et al.\\ (1993) and Abraham \\& Carrara (1998) claim that these can be explained by ejection of components by a precessing jet.  3C~279 was observed during the TDRSS space VLBI experiments at 2.3\\,GHz (Linfield et al.\\ 1989) and 15 GHz (Linfield et al.\\ 1990), with source frame brightness temperatures between 1.6 and 2.0~$\\times$~10$^{12}$~K being measured for this source. A brightness temperature of 1.9~$\\times$~10$^{12}$~K (translated from observed frame to source frame) was also measured for 3C~279 in the 22~GHz VSOP Pre-Launch Survey (Moellenbrock et al.\\ 1996). This source has been detected in VLBI observations up to frequencies of 215~GHz (Krichbaum et al.\\ 1997). This paper reports on the VSOP observations of 3C~279 made during the first Announcement of Opportunity period (AO1). Hirabayashi et al.\\ (1999) and Edwards et al.\\ (1999) presented preliminary analyses of these 5 and 1.6 GHz VSOP observations respectively.  Here we present higher dynamic range images together with analysis and interpretation of model fits and a spectral index map. ",
        "conclusions": "We have presented the first space VLBI images of 3C 279, which are the highest resolution images yet obtained of this source at 5 and 1.6 GHz.  The parsec-scale emission is dominated by the core and the jet component C4 which has been visible in VLBI images since 1985. The 1.6 GHz VSOP image and the 5 and 1.6 GHz VLBA-only images show emission from a jet extending to the southwest.  The 1.6 GHz VLBA-only image has structure that matches that seen in the highest resolution VLA images, connecting the parsec and kiloparsec scale structures in this source.  We have exploited two of the main strengths of VSOP: the ability to obtain matched-resolution images to ground-based images at higher frequencies and the ability to measure high brightness temperatures.  The spectral index map constructed from the 1.6 GHz VSOP image and the 5 GHz VLBA-only image has an unusually inverted spectral index in the core region, approaching the limiting value for synchrotron self-absorption of $+2.5$. An extensive error analysis conducted on the model-fit brightness temperatures reveals brightness temperature lower limits as high as $5~\\times~10^{12}$~K. This lower limit is significantly above both the nominal inverse Compton and equipartition brightness temperature limits.  The derived Doppler factor, Lorentz factor, and angle to the line-of-sight in the case of the equipartition limit are at the upper end of the range of expected values for EGRET blazars."
    },
    "0002/astro-ph0002020_arXiv.txt": {
        "abstract": "We present the results of fitting deep off-nuclear optical spectra of radio-quiet quasars, radio-loud quasars and radio galaxies at $z \\simeq 0.2$ with evolutionary synthesis models of galaxy evolution. Our aim was to determine the age of the dynamically dominant stellar populations in the host galaxies of these three classes of powerful AGN. Some of our spectra display residual nuclear contamination at the shortest wavelengths, but the detailed quality of the fits longward of the 4000\\AA\\ break provide unequivocal proof, if further proof were needed, that quasars lie in massive galaxies with (at least at $z \\simeq 0.2$) evolved stellar populations. By fitting a two-component model we have separated the very blue (starburst and/or AGN contamination) from the redder underlying spectral energy distribution, and find that the hosts of all three classes of AGN are dominated by old stars of age $8 - 14$ Gyr. If the blue component is attributed to young stars, we find that, at most, 1\\% of the visible baryonic mass of these galaxies is involved in star-formation activity at the epoch of observation, at least over the region sampled by our spectroscopic observations. These results strongly support the conclusion reached by McLure et al. (1999) that the host galaxies of luminous quasars are massive ellipticals which have formed by the epoch of peak quasar activity at $z \\simeq 2.5$. ",
        "introduction": "Determining the nature of the host galaxies of powerful active galactic nuclei (AGN) is of importance not only for improving our understanding of different manifestations of AGN activity, but also for exploring possible relationships between nuclear activity and the evolution of massive galaxies. The recent affirmation that black-hole mass appears approximately proportional to spheroid mass in nearby inactive glaxies (Maggorian et al. 1998) has further strengthened the motivation for exploring the link between active AGN and the dynamical and spectral properties of their hosts.  Over the last few years, improvements in imaging resolution offered by both space-based and ground-based optical/infrared telescopes has stimulated a great deal of research activity aimed at determining the basic structural parameters ({\\it i.e.} luminosity, size and morphological type) of the hosts of radio-loud quasars, radio-quiet quasars, and lower-luminosity X-ray-selected and optically-selected AGN ({\\it e.g.} Disney et al. 1995; Hutchings \\& Morris 1995; Bahcall et al. 1997; Hooper et al. 1997; McLure et al. 1999, Schade et al. 2000). However, relatively little corresponding effort has been invested in spectroscopic investigations of AGN hosts, despite the fact that this offers an independent way of classifying these galaxies, as well as a means of estimating the age of their stellar populations.  Our own work in this field to date has focussed on the investigation of the hosts of matched samples of radio-quiet quasars (RQQs), radio-loud quasars (RLQs) and radio galaxies (RGs) at relatively modest redshift ($z = 0.2$). Details of these samples can be found in Dunlop et al. (1993). In brief, the sub-samples of quasars (i.e. RQQs and RLQs) have been selected to be indistinguishable in terms of their two-dimensional distribution on the $V-z$ plane, while the sub-samples of radio-loud AGN (i.e. RLQs and RGs) have been selected to be indistinguishable in terms of their two-dimensional distribution on the $P_{5GHz} - z$ plane (as well as having indistinguishable spectral-index distributions). Deep infrared imaging of these samples (Dunlop et al. 1993; Taylor et al. 1996) has recently been complemented by deep WFPC2 HST optical imaging (McLure et al. 1999), the final results of which are reported by Dunlop et al. (2000). As well as demonstrating that, dynamically, the hosts of all three types of luminous AGN appear indistinguishable from normal ellipticals, this work has enabled us to deduce crude spectral information on the host galaxies in the form of optical-infrared colours. However, broad-baseline colour information can clearly be most powerfully exploited if combined with detailed optical spectroscopy. Over the past few years we have therefore attempted to complement our imaging studies with a programme of deep optical off-nuclear spectroscopy of this same sample of AGN.  Details of the observed samples, spectroscopic observations, and the basic properties of the observed off-nuclear spectra are given in a companion paper (Hughes et al. 2000). As discussed by Hughes et al. (2000), the key feature of this study (other than its size, depth, and sample control) is that we have endeavoured to obtain spectra from positions further off-nucleus ($\\simeq 5$ arcsec) than previous workers, in an attempt to better minimize the need for accurate subtraction of contaminating nuclear light. This approach was made possible by our deep infrared imaging data, which allowed us to select slit positions $\\simeq 5$ arcsec off-nucleus which still intercepted the brighter isophotes of the host galaxies (slit positions are shown, superimposed on the infrared images, in Hughes et al. 2000).   In this paper we present the results of attempting to fit the resulting  off-nuclear spectra with evolutionary synthesis models of galaxy evolution. Our primary aim was to determine whether, in each host galaxy, the optical spectrum could be explained by the same model as the optical-infrared colour, and, if so, to derive an estimate of the age of the dynamically dominant stellar population. However, this study also offered the prospect of determining whether the hosts of different classes of AGN differ in terms of their more recent star-formation activity.  It is worth noting that our off-nuclear spectroscopy obviously does not enable us to say anything about the level of star-formation activity in the nucleus of a given host galaxy. Rather, any derived estimates of the level of star-formation activity refer to the region probed by our observations $\\simeq 5$ arcsec off-nucleus. However, given the large scalelengths of the hosts, their relatively modest redshift, and the fact that our spectra are derived from long-slit observations, it is reasonable to regard our conclusions as applying to fairly typical regions, still located well within the bulk of the host galaxies under investigation.  The layout of this paper is as follows. In section 2 we provide details of the adopted models, and how they have been fitted to the data. The results are presented in section 3, along with detailed notes on the fitting of individual spectra. The implications of the model fits are then discussed in section 4, focussing on a comparison of the typical derived host-galaxy ages in the three AGN subsamples. Finally our conclusions are summarized in section 5. The detailed model fits, along with corresponding chi-squared plots are presented in Appendix A. ",
        "conclusions": "We conclude that the hosts of all three major classes of AGN contain predominantly old stellar populations ($\\simeq 11$ Gyr) by $z \\simeq 0.2$. This agrees well with the results of McLure \\etal (1999), and Dunlop \\etal (2000) who compare host galaxy morphologies, luminosities, scale lengths and colours in the same sample, and conclude that the hosts are, to first order, indistinguishable from `normal' quiescent giant elliptical galaxies.  The best-fitting age of the dominant stellar population is {\\it not} a function of AGN class. For the purely stellar models, the fitted percentage contribution of the blue component is, however, greater in the quasar hosts than in the radio galaxies; the median values are 0.6\\% for the 9 radio-loud quasars, 0.6\\% for the 9 radio-quiet quasars, and 0.05\\% for the 6 radio galaxies. However, when a nuclear component is included, the median values are 0.3\\% for the radio-loud quasars, 0.3\\% for the radio-quiet quasars, and 0.00\\% for the radio galaxies. Performing a Kolmogorov-Smirnov test on these results yields a probability greater than 0.2 that the percentage of young stellar population in host galaxies is in fact also not a function of AGN class.   These results strongly support the conclusion that the host galaxies of all three major classes of AGN are massive ellipticals, dominated by old stellar populations. \\\\ \\\\ \\\\ {\\bf ACKNOWLEDGEMENTS}\\\\ \\\\ Louisa Nolan acknowledges the support provided by the award of a PPARC Studentship. Marek Kukula and David Hughes acknowledge the support provided by the award PPARC PDRAs, while Raul Jimenez acknowledges the award of a PPARC Advanced Fellowship. We thank an anonymous referee for perceptive comments which helped to clarify the robustness of our results and improved the clarity of the paper. \\\\ {\\bf REFERENCES}\\\\ \\\\ Bahcall J.N., Kirhakos S., Saxe D.H., Schneider D.P., 1997, ApJ, 479, 642\\\\ Bruzual A.G., Charlot S., 2000, in preparation\\\\ Disney M.J. et al., 1995, Nat, 376, 150\\\\ Dunlop J.S., 2000, In: `The Hy-Redshift Universe: Galaxy Formation and Evolution at High Redshift', ASP Conf. Ser., Vol 193, eds. A.J. Bunker \\& W.J.M. van Breugel, in press, (astro-ph/9912380)\\\\ Dunlop J.S., McLure R.J., Kukula M.J., Baum S.A., O'Dea C.P., Hughes D.H., 2000, MNRAS, in press\\\\ Dunlop J.S., Taylor G.L., Hughes D.H., Robson E.I., 1993, MNRAS, 264, 455\\\\ Hooper E.J., Impey C.D., Foltz C.B., 1997, ApJ, 480, L95\\\\ Hughes D. H., Kukula M.J., Dunlop J.S., Boroson T., 2000, MNRAS, in press\\\\ Hutchings J.B., Morris S.C., 1995, AJ, 109, 1541\\\\ Jimenez R., Dunlop J.S., Peacock J.A., Padoan P., MacDonald J., J$\\o$rgensen U.G., 2000, MNRAS, in press\\\\ Jimenez R., Padoan P., Matteucci F., Heavens A.F., 1998, MNRAS, 299, 123\\\\ Magorrian J., et al., 1998, AJ, 115, 2285\\\\ McLure R.J., Kukula M.J., Dunlop J.S., Baum S.A., O'Dea C.P., Hughes D.H., 1999, MNRAS, in press, (astro-ph/9809030)\\\\ Nolan L.A., Dunlop J.S., Jimenez R., 2000, MNRAS, submitted, (astro-ph/0004325) Schade, D.J., Boyle, B.J., Letawsky, M., 2000, MNRAS, in press\\\\ Taylor G.T., Dunlop J.S., Hughes D.H., Robson E.I., 1996, MNRAS, 283, 930\\\\ Yi S., Brown T.M., Heap S., Hubeny I., Landsman W., Lanz T., Sweigart A., 2000, ApJ, submitted\\\\    \\appendix"
    },
    "0002/astro-ph0002335_arXiv.txt": {
        "abstract": "This paper presents infrared spectra taken with the newly commissioned NIRSPEC spectrograph on the Keck II Telescope of the High Redshift Radio Galaxy MRC~2025-218 (z=2.63) These observations represent the deepest infrared spectra of a radio galaxy to date and have allowed for the detection of H$\\beta$, [OIII] (4959/5007), [OI] (6300), H$\\alpha$, [NII] (6548/6583) and [SII] (6716/6713).  The H$\\alpha$ emission is very broad (FWHM = 9300 km/s) and luminous (2.6$\\times$10$^{44}$ ergs/s) and it is very comparable to the line widths and strengths of radio loud quasars at the same redshift. This strongly supports AGN unification models linking radio galaxies and quasars, although we discuss some of the outstanding differences.  The [OIII] (5007) line is extremely strong and has extended emission with large relative velocities to the nucleus.  We also derive that if the extended emission is due to star formation, each knot has a star formation rate comparable to a Lyman Break Galaxy at the same redshift. ",
        "introduction": "Deep radio surveys have proven to be one of the best methods for finding high redshift galaxies. Most evidence suggests that these powerful radio sources are the precursors of local giant ellipticals (e.g. Pentericci, et al. 1999).  Many have irregular and complex morphologies suggestive of mergers and they are often surrounded by an overdensity of compact sources; presumably sub-galactic clumps (e.g. van Breugel et al. 1998).  At both low and high redshifts, radio galaxies usually have strong optical emission lines, especially OIII at 5007~\\AA.  It is strongly debated, however, if the emission lines arise by the same mechanism as the radio jets. Several authors (e.g. Rawlings \\& Saunders 1992, Eales \\& Rawlings 1993, and Evans 1998) have demonstated a strong correlation between radio luminosity and [OIII] luminosity, but as Evans showed, there is a strong selection effect based on the detection limits as a function of distance and this may explain much of the correlation.  Since the galaxies are often disturbed, star formation, large scale shocks and a central AGN are all possible sources of the line emission.  Active galaxy unification models suggest that radio galaxies are quasars with obscured broad line regions (e.g. Antonucci 1993). Eales \\& Rawlings (1993, 1996) and Evans (1998) have been successful at using infrared spectrographs on 4-meter class telescopes to measure a few of the brightest lines in small samples of radio galaxies in the redshift range 2.2 to 2.6 and they find line ratios most consistent with Seyfert 2 (obscured AGN) nuclei.  Independent of our current efforts, a team has also successfully used the ISAAC instrument on the VLT to observe high redshift radio galaxies (HzRGs) including MRC~2025-218 (McCarthy, personal communication).  An additional unexplained phenomenon is that at high redshifts (z$>$0.6) the radio, optical continuum, infrared continuum and emission line structures tend to be closely aligned (Chambers, Miley, \\& van Bruegel 1987, and McCarthy et al. 1987). This is probably not seen in lower redshift targets because the central activity tends to be a smaller fraction of the total luminosity than in high redshift sources. Among the proposed explanations are that the emission lines arise from shock induced star formation (De Young 1989, Rees 1989) or that it is scattered light originating from the central nucleus (Fabian 1991).  This is a crucial question in our understanding of how and when most star formation occured in giant elliptical galaxies and in clusters in general.  MRC~2025-218 (z=2.630) has a compact near infrared and optical continuum morphology (van Breugel et al. 1998), but extended Ly$\\alpha$ emission (5$^{\\prime \\prime}$) aligned with its radio axis (McCarthy et al. 1992). The extended UV emission has significant (8.3$\\pm$2.3 \\%) linear polarization perpendicular to the UV axis (Cimatti et al. 1996) suggesting that scattering plays a significant role.  McCarthy et al. also found three extremely red galaxies (ERO's: R-K $>$ 6 mag) within 20$^{\\prime \\prime}$ of the radio galaxy.  This is a large overdensity of such objects and strongly suggests an association between the ERO's and the active galaxy.  In this paper we present infrared spectra taken with a long slit oriented close to the radio axis and including one of the ERO's.  The ERO spectra will be described in a future paper. For all calculations we have assumed a cosmology with $\\Lambda$=0, q$_0$=0.1 and H$_0$=75 km s$^{-1}$ Mpc$^{-1}$.  For a redshift of 2.63 this yields a luminosity distance of 2.1$\\times$10$^4$ Mpc and an angular scale of 7.7 kpc per arcsecond. ",
        "conclusions": "We have obtained the most sensitive infrared spectra ever taken of a high redshift radio galaxy. The galaxy has very strong emission lines with ratios and line widths consistent with an obscured quasar.  The narrow line region appears to be partially obscured with A$_V$ around 1.4 mag, but from comparisons with high redshift quasars, we estimate that the extinction to the broad line region is between 3 and 5 magnitudes.  Since other radio galaxies in the same redshift range don't show broad emission lines, we suggest that MRC~2025-218 is further along in its evolution towards an unobscured quasar.  We cannot rule out any of the proposed mechanisms for the production of the aligned emission. But based on the [OIII] line strength if the majority of the emission is due to star formation, we find that the star formation rate is comparable to that of Lyman Break Galaxies at similar redshifts.  We urge even deeper observations of this and other similar radio galaxies in order to measure additional extended line emission."
    },
    "0002/astro-ph0002273_arXiv.txt": {
        "abstract": "We have analysed the ISO-SWS spectrum of Jupiter in the 12--16 $\\mu{\\rm m}$ range, where several hydrocarbons exhibit rovibrational bands. Using temperature information from the methane and hydrogen emissions, we derive the mixing ratios ($q$) of acetylene and ethane at two independent pressure levels. For acetylene, we find $q=(8.9^{+1.1}_{-0.6})\\times10^{-7}$ at 0.3 mbar and $q=(1.1^{+0.2}_{-0.1})\\times10^{-7}$ at 4 mbar, giving a slope $-d\\ln q / d\\ln P=0.8\\pm0.1$, while for ethane $q=(1.0\\pm0.2)\\times10^{-5}$ at 1 mbar and $q=(2.6^{+0.5}_{-0.6})\\times10^{-6}$ at 10 mbar, giving $-d\\ln q / d\\ln P=0.6\\pm0.2$. The ethane slope is consistent with the predictions of Gladstone et al.\\ (\\cite{Gladstone96}), but that predicted for acetylene is larger than we observe. This disagreement is best explained by an overestimation of the acetylene production rate compared to that of ethane in the Gladstone et al.\\ (\\cite{Gladstone96}) model.  At 15.8 $\\mu{\\rm m}$, methylacetylene is detected for the first time at low jovian latitudes, and a stratospheric column density of $(1.5\\pm0.4)\\times10^{15}$ molecule\\,cm$^{-2}$ is inferred. We also derive an upper limit for the diacetylene column density of $7\\times10^{13}$ molecule\\,cm$^{-2}$. ",
        "introduction": "\\indent\\indent Hydrocarbons in Jupiter are produced in a series of chemical pathways initiated by the photolysis of methane in the upper stratosphere. Vertical transport, mainly turbulent diffusion, redistributes the molecules throughout the stratosphere and down to the troposphere, where they are eventually destroyed. Hydrocarbons, and particularly the most stable of them, are therefore good tracers of the upper atmospheric dynamics. In addition, as they strongly contribute to the atmospheric opacity in the UV and IR, hydrocarbons act as major sources and sinks of heat, thereby participating to the stratospheric dynamics.  All these reasons have strongly motivated theoretical studies of the jovian stratospheric photochemistry (Strobel \\cite{Strobel69}; Yung \\& Strobel \\cite{Yung80}; and most recently Gladstone et al.\\ \\cite{Gladstone96}). Although nowadays very detailed, these models still need to be constrained by observations of minor species. However, prior to the ISO mission, only two molecules, acetylene (${\\rm C_2H_2}$) and ethane (${\\rm C_2H_6}$), had been detected, except in the auroral zones where several other minor species (${\\rm C_2H_4}$, ${\\rm C_3H_4}$ and ${\\rm C_6H_6}$) have been observed (Kim et al.\\ \\cite{Kim85}). Mean stratospheric mole fractions have been inferred for ${\\rm C_2H_2}$ and ${\\rm C_2H_6}$ by various authors, but no precise information on their vertical variations was made available. In this paper, we analyse the ISO-SWS spectrum of Jupiter between 7 and 17 $\\mu{\\rm m}$, in order to determine the vertical distributions of ${\\rm C_2H_2}$ and ${\\rm C_2H_6}$, and to search for more complex (C$_3$ and C$_4$) molecules. Sect.~2 presents the observations. Our analysis of the spectrum is presented in Sect.~3. The results are compared with previous observations and theoretical predictions in Sect.~4. ",
        "conclusions": "\\indent\\indent B\\'ezard et al.\\ (\\cite{Bezard95}) first showed from 13.4-$\\mu{\\rm m}$ high-resolution spectroscopy that the ${\\rm C_2H_2}$ mixing ratio increases with height in the stratosphere, and that most of the acetylene is concentrated above the $\\sim$0.5-mbar level. Their distribution has a mixing ratio of about $7\\times10^{-7}$ at 0.3 mbar, in reasonable agreement with our results. More recently, B\\'etremieux \\& Yelle (\\cite{Betremieux99}), using UV observations, found an average ${\\rm C_2H_2}$ mixing ratio in the 20--60 mbar range of $1.5\\times10^{-8}$, consistent with our value of $1.9\\times10^{-8}$, extrapolated to this pressure range. Using height-dependent mixing ratio profiles to analyse high-resolution infrared observations, Sada et al.\\ (\\cite{Sada98}) derived mixing ratios of $3.9^{+1.9}_{-1.3}\\times10^{-6}$ for ${\\rm C_2H_6}$ at 5 mbar and $2.3\\pm0.5\\times10^{-8}$ for ${\\rm C_2H_2}$ at 8 mbar. While the former value exactly agrees with our results, the latter is almost 3 times less than that extrapolated downwards from our ${\\rm C_2H_2}$ profile. Also from infrared observations, Noll et al.\\ (\\cite{Noll86}), using a slope of $-d\\ln q / d\\ln P=0.3$, found a ${\\rm C_2H_6}$ mixing ratio of $7.5\\times10^{-6}$ at 1 mbar, which compares reasonably well with our measurement at the same pressure level.  We also compared our retrieved mixing ratios with those calculated in the photochemical model of Gladstone et al.\\ (\\cite{Gladstone96}). For ${\\rm C_2H_6}$, our results are in excellent agreement with their model both at 1 and 10 mbar (Fig.~\\ref{FigPro}). For ${\\rm C_2H_2}$, while the agreement is good at 4 mbar, their mixing ratio at 0.3 mbar is higher than ours by a factor of 4. Our ${\\rm C_2H_2}$ slope is then significantly lower than theirs ($-d\\ln q / d\\ln P=0.8\\pm0.1$ vs.\\ $-d\\ln q / d\\ln P=1.2$). Note that our derived ${\\rm C_2H_2}$ slope is still steeper than that of ${\\rm C_2H_6}$ ($-d\\ln q / d\\ln P=0.6\\pm0.2$), as expected. Indeed, ${\\rm C_2H_6}$, being less subject to photolysis, has a longer lifetime in the jovian stratosphere than ${\\rm C_2H_2}$ (Gladstone et al.\\ \\cite{Gladstone96}, their Fig.~6).  Hydrocarbons are formed from the photolysis of ${\\rm CH_4}$ around the homopause. Small-scale turbulence, parameterized in a photochemical model by the eddy diffusion coefficient ($K$), transports them downwards in the stratosphere. This process is approximately modelled for long-lived products by the equation $K(z)n(z){\\rm d}q/{\\rm d}z = P(z)$, where $n(z)$ is the number density at altitude $z$ and $P(z)$ the vertically integrated net production rate above altitude $z$. A first hypothesis would attribute the difference between the observed and calculated ${\\rm C_2H_2}$ abundances to an overestimation of the ${\\rm C_2H_2}$ production rate $P(z)$ in the Gladstone et al.\\ model. Our two abundance measurements at 4 and 0.3 mbar allow us to calculate the mean ${\\rm d}q/{\\rm d}z$ over this pressure range. Comparing with the value of ${\\rm d}q/{\\rm d}z$ predicted by Gladstone et al.\\ at 1 mbar, we found that $P(z)$ should be decreased by a factor of $\\sim$3.  A second explanation would be that Gladstone et al.\\ underestimated the eddy diffusion coefficient $K$ by a factor of $\\sim$3 around 1 mbar. This underestimation at 1 mbar should also apply to the level and eddy diffusion coefficient ($K_H$) at the homopause. However, it is difficult to directly quantify the changes on the homopause parameters, because of the strong coupling between production rates and homopause level. Simply, we note that our analysis could imply an increase in $K_H$. It is consistent with Drossart et al.\\ (\\cite{Drossart99}), who, analysing ${\\rm CH_4}$ fluorescence, found $K_H=(7\\pm1)\\times10^6$ ${\\rm cm^{-2}\\,s^{-1}}$, higher than the value of 1.4$\\times10^6$ assumed in the Gladstone et al.\\ model. In this case, the agreement on the ${\\rm C_2H_6}$ slope would also imply that the ${\\rm C_2H_6}$ production rate has been underestimated by Gladstone et al.\\ (\\cite{Gladstone96}). In fact, the most direct conclusion of our measurements is that Gladstone et al.\\ have overestimated the ${\\rm C_2H_2}$/${\\rm C_2H_6}$ production rate ratio.  The ISO-SWS spectrum enables the first detection of ${\\rm CH_3C_2H}$ in the equatorial region of Jupiter. We retrieved a column density of $(1.5\\pm0.4)\\times10^{15}$ molecule\\,cm$^{-2}$. Kim et al.\\ (\\cite{Kim85}) had previously detected ${\\rm CH_3C_2H}$ in the north polar auroral zone of Jupiter, and had retrieved a column density of $(2.8^{+2.4}_{-1.1})\\times10^{16}$ molecule\\,cm$^{-2}$. At least a large part of the difference is explained by different modelling assumptions. Kim et al.\\ (\\cite{Kim85}) assumed a uniform vertical distribution throughout the stratosphere and used a temperature profile for the auroral region which is now known to be incorrect (Drossart et al.\\ \\cite{Drossart93}). For ${\\rm C_4H_2}$, we found only an upper limit of $7\\times10^{13}$ molecule\\,cm$^{-2}$, 65 times lower than the stratospheric column density predicted by Gladstone et al.\\ (4.5$\\times$10$^{15}$ molecule\\,cm$^{-2}$). This is consistent with their overestimate of ${\\rm C_2H_2}$ production, since the production of ${\\rm C_4H_2}$ is essentially quadratically dependent on the abundance of ${\\rm C_2H_2}$. As our ${\\rm CH_3C_2H}$ column density is 3.5 times larger than calculated by Gladstone et al., we derived a ${\\rm CH_3C_2H}$/${\\rm C_4H_2}$ ratio larger than 20, while they found a ratio of $\\sim2$. This discrepancy is all the more remarkable as, in the case of Saturn, where both ${\\rm CH_3C_2H}$ and ${\\rm C_4H_2}$ are detected (de Graauw et al.\\ \\cite{deGraauw97}), the photochemical model of Moses et al.\\ (\\cite{Moses99}) gives, as observed, a ${\\rm CH_3C_2H}$/${\\rm C_4H_2}$ ratio of about 10. This stresses that the C$_3$ and C$_4$ chemistry in Jupiter should be reassessed with a complete photochemical model of the jovian stratosphere."
    },
    "0002/astro-ph0002209_arXiv.txt": {
        "abstract": "We present spectrophotometric imaging mid--IR observations of the reflection nebula Ced\\,201. Ced\\,201 is a part of a molecular cloud illuminated by a B9.5V star moving through it at more than 12 km s$^{-1}$.  The spectra of Ced\\,201 give evidence for transformation  of very small carbonaceous grains into the carriers of the Aromatic Infrared Bands (AIBs), due to the radiation field of the illuminating star and/or to shock waves created by its motion. These very small grains emit mainly very broad bands and a continuum. We suggest that they are present everywhere in the interstellar medium but can only be detected in the mid--IR under special circumstances such as those prevailing in this reflection nebula. The efficiency of energy conversion of stellar light into mid--infrared emission is 7.5\\%~  for both the very small grains and the AIB carriers, and the fraction of interstellar carbon locked in these emitters is approximately 15\\%. \\keywords\t{ISM: Ced\\,201\t\t\t- dust, extinction\t\t- Infrared: ISM: lines and bands} ",
        "introduction": "The mid--infrared emission bands at 3.3, 6.2, 7.7, 8.6, 11.3 and 12.7\\,\\mum~ are ubiquitous in the interstellar medium (ISM). The coincidence of these bands with characteristic wavelengths of aromatic molecules and attached functional groups (see e.g. Allamandola et al. \\cite{Allamandola89}) is sufficiently convincing to call them Aromatic Infrared Bands (AIBs). Other, more  neutral designations are Unidentified Infrared Bands (UIBs) or Infrared  Emission Features (IEFs). The designation should stabilize in the future.  Their carriers are very small grains (or big molecules) heated transiently by the absorption of a single  UV (Sellgren et al. \\cite{Sellgren83}) or visible (Uchida et al.  \\cite{Uchida98};  Pagani et al. \\cite{Pagani99}) photon. They are often identified  with Polycyclic Aromatic Hydrocarbons (PAHs) (L\\'eger \\& Puget \\cite{Leger}, Allamandola et al. \\cite{Allamandola85}). Duley \\& Williams (\\cite{Duley88}) and Duley et al. (\\cite{Duley93}) have proposed an alternative interpretation  in which the carriers of the AIBs are not free particles but ``graphitic  islands'' (i.e. akin to PAHs) loosely attached to each other so that the heat  conductivity is small. Then their thermal behaviour is comparable to that of  free particles. Similar models have been also proposed by others. The carriers could also be \\mbox{3--D} very small particles. In this Letter, we will not enter  into this debate and will simply call them the {\\it AIB carriers}.  The origin of the AIB carriers is still unclear. The idea that they are formed in the envelopes of carbon stars is not supported by observation and meets with theoretical difficulties (e.g. Cherchneff et al. \\cite{Cherchneff}).  They might be formed outside carbon star envelopes from  carbonaceous particles condensed previously in these envelopes. Schnaiter et al. (\\cite{Schnaiter99}) and others have proposed such an evolutionary scheme from carbon stars to planetary nebulae. However the AIB carriers formed in this way might not survive the strong UV field of planetary nebulae. For example, Cox et al. (\\cite{Cox}) find no AIB in the Helix nebula, an old carbon--rich planetary nebula. Another possibility is that they are formed in the ISM from carbonaceous grains or carbonaceous mantles covering silicate nuclei. Boulanger et al. (\\cite{Boulanger}) and Gry et al. (\\cite{Gry}) have suggested that the strong variations in the IRAS 12\\,\\mum/100\\,\\mum~ interstellar flux ratio are due to the release of transiently heated particles (emitting at 12\\,\\mum) by bigger particles, respectively due to  UV irradiation or to shattering by shocks. Thanks to the Infrared Space Observatory (ISO), we know now that in the general ISM the emission in the  IRAS 12\\,\\mum~ band is dominated by AIBs. Variations in the ratio of AIB carriers to big grains are thus confirmed. Laboratory experiments shed  some light  on this process. Scott et al. (\\cite{Scott97}) have observed the release of aromatic carbon clusters containing in excess of 30 carbon atoms  by solid hydrogenated amorphous carbon (HAC)  irradiated by a strong UV laser pulse. The energy deposited by the laser pulse per unit target area is similar to the energy of a central  collision between two grains of 10 nm radius at a velocity of 10 km s$^{-1}$. This is an indication for a possible release of AIB carriers in interstellar  shocks.  ISO has also shown the existence of variations of AIB spectra that may be related to a transformation of carbonaceous grains.  Recently, Uchida et al. (\\cite{Uchida00}) have observed a broadening of the 6.2, 7.7 and 8.6\\,\\mum~ AIBs when going from strong to weak radiation field in the reflection nebula vdB\\,17 = NGC\\,1333. They also notice an increase of the flux ratio $I$(5.50--9.75\\,\\mum)/$I$(10.25--14.0\\,\\mum) with increasing radiation field in several reflection nebulae.  We present mid--IR spectrophotometric and CO(2--1) line observations of another reflection nebula, Ced\\,201, that shed light into this transformation. Section 2 describes briefly the observations and reductions. Sect. 3 describes the results, and Sect. 4  contains a discussion and the conclusions. ",
        "conclusions": "Fig. 3 shows a behaviour of the spectrum with radiation field similar to that observed by Uchida et al. (\\cite{Uchida00}) in some other reflection nebulae.  We find a trend for the 7.7\\,\\mum~AIBs to become broader at {\\it fainter} radiation fields away from the exciting star, Fig. 4. The trend is not clear for the other bands, but we should not forget that there is contamination by the AIBs from foreground and background material. We suggests that some carbonaceous material that emits the continuum and broad bands far from the star is processed through the effect of the star that moves through the molecular cloud, producing AIB carriers: the continuum is only 2 times fainter 12\\arcsec~ from the star than close to the star, demonstrating the partial disappearance of its carriers near the star while the AIBs become very strong. The very appearance at relatively large distances from the star (at least 18\\arcsec) of a continuum rising towards long wavelengths is surprising.   This continuum must be emitted by VSGs heated by single (visible) photons. These grains must be quite smaller than the ``classical'' VSGs which require a strong radiation field to emit in the wavelength range of ISOCAM (Contursi et al. \\cite{Contursi}). The spectra seen far from the star remind strongly of Class B spectra seen in carbon--rich proto--planetary nebulae (Guillois et al. \\cite{Guillois96}). In the laboratory, such spectra are produced (in absorption) by natural coals rich in aromatic cycles (Guillois et al. \\cite{Guillois96}) or by a-C:H materials produced by laser pyrolysis of hydrocarbons (Herlin et al. \\cite{Herlin98}; Schnaiter et al. \\cite{Schnaiter99}). In Ced\\,201, these carbonaceous grains must be very small (radius of the order of 1 nm) since they are heated transiently by visible photons to the temperatures of $\\simeq$ 250 K necessary to emit the observed mid--IR features. Since this is also the approximate size of the classical AIB carriers it would not be appropriate to say that these particles release these carriers. One should invoke instead chemi--physical transformations like aromatization (Ryter \\cite{Ryter}). In any case, these small grains were already present before the star penetrated the molecular cloud. They are expected to be present elsewhere in the ISM.   \\begin{figure}[t!] \\vspace{-0.6cm} \\psfig{file=ced_width.ps,width=8.8cm,height=8.8cm,angle=0.0} \\caption{Lorentzian widths of four AIB features from Fig. 3 plotted as a function of the angular distance to the exciting star, normalized to the width of the Lorentzian at zero distance. The error bars show the internal dispersion, i.e. the dispersion of values obtained at the same distance from the central star (i.e. 1 value at distance 0 -- hence null dispersion, 4 values at distance 6 arcsec, etc.)} \\end{figure}  We find that assuming spherical symmetry the total emissivity per unit  volume in the CVF spectral range (5.15 to 17\\,\\mum) decreases approximately  as the inverse square root of the distance to the star, independently of the shape of the spectrum. We base the following calculation on the light diffusion model of Witt et al.  (\\cite{Witt}), who postulate a uniform density that we find equal to n(H) = 2n(H$_2$) = 1800 cm$^{-3}$. This density corresponds to a visual extinction of only 0.02 mag. per 6\\arcsec~ angular distance. Since the integrated mid--IR emission is proportional to the stellar  radiation density, we can derive an energy conversion efficiency, $\\eta$ = 7.5 \\%~ (emitted in 4$\\pi$ steradians). Knowing the density, the  radiation field of the B9.5\\,V star, and adopting an absorption cross--section  per carbon atom of $3\\times 10^{-18}$ cm$^{-2}$ (Allamandola et al.  \\cite{Allamandola89}) and an interstellar carbon abundance C/H=$2.3\\times 10^{-4}$ (Snow \\& Witt \\cite{Snow95};  see also Andrievsky et al. \\cite{Andrievsky} for carbon abundance in B stars), we find that the fraction of carbon locked in the very small particles emitting in the mid--IR is 15\\% . This is in agreement with the independent analysis of Kemper et al. (\\cite{Kemper}) who find that 10\\%  of the gas--phase carbon is locked in particles  smaller than 1.5 nm. The {\\it absorption} by these particles contributes  to 1/4 of the {\\it extinction} in the visible, a considerable fraction.  We have examined the ratio of the strength of the 6.2\\mum~ (a vibration of the aromatic skeleton) and of the 11.3\\mum~AIB  features (a C--H bending mode) as a function of distance to the exciting star. No variation has been seen in this ratio within our rather large (30 \\%) errors.  It is interesting to note that Witt et al. (\\cite{Witt}) find from UV--visible scattering studies of Ced\\,201 that the grains responsible for the visible scattering have a narrow size distribution skewed towards big, wavelength--sized grains. This is mostly seen within about 20\\arcsec~ from the star. On the other hand the very strong 2175 \\AA~ extinction band suggests the presence of many smaller carbon grains. All this is evidence for grain processing by the radiation of the star, or/and by shattering of grains by a shock wave associated with its motion. The very small carbonaceous grains we see in Ced\\,201 through their mid--IR emission might be responsible for the 2175 \\AA~ band.  In summary, we have shown the existence of transformations near a star of very small, \\mbox{3--D} hydrogenated carbonaceous grains which are probably  present everywhere in the ISM. This transformation of carbonaceous grains into AIB carriers might be due either to the radiation of a B9.5V star, or to a shock induced by the supersonic  motion of the star through the molecular cloud. Although the  existence of a shock is indicated by an optical arc visible 18\\arcsec~ from the star (and also suspected from our CO observations), there is no  obvious sign of it in the spectra shown in Fig. 3 and it may not be the cause of the transformation. The very small grains as well as the AIB carriers formed  from them are excited by the light of the star. The conversion efficiency between received excitation energy and mid--IR emitted energy is of the order of 7.5\\%  for both kinds of particles, and the fraction of carbon locked in these particles is approximately 15\\%."
    },
    "0002/astro-ph0002515_arXiv.txt": {
        "abstract": "We describe the nature of the interference challenges facing radio astronomy in the next decade. These challenges will not be solved by regulation only, negotiation and mitigation will become vital. There is no silver bullet for mitigating against interference. A successful mitigation approach is most likely to be a hierarchical or progressive approach throughout the telescope and signal conditioning and processing systems. We summarise some of the approaches, including adaptive systems. ",
        "introduction": "If future telescopes like the SKA are developed with sensitivities up to 100 times greater than present sensitivities, it is quite likely that current regulations will not provide the necessary protection against interference. There is a range of experiments (eg redshifted hydrogen or molecular lines) which require use of arbitrary parts of the spectrum, but only at a few locations, and at particular times, suggesting that a very flexible approach may be beneficial. Other experiments require very large bandwidths, in order to achieve enough sensitivity. As shown in Figure \\ref{alloc}, presently only 1-2\\% of the spectrum in the metre and centimetre bands is reserved for passive uses, such as radio astronomy (Morimoto 1993). In the millimetre band, much larger pieces of the spectrum are available for passive use, but the existing allocations are not necessarily at the most useful frequencies. Current regulations alone will be inadequate, we need technology as well as regulation.  We cannot (and do not want to) impede the telecommunications revolution, but we can try to minimise its impact on passive users of the radio spectrum and maximise the benefits of technological advances. Further information on many of the topics discussed below is available on http://www.atnf.csiro.au/SKA/intmit/.    \\begin{figure}[htbp] \\begin{center} \\plotone{spectrum_allocation2.ps} \\caption{Spectral lines (at zero redshift) are indicated in absorption or emission over the 0--30 GHz band. The boxes indicate the bands allocated for passive radio astronomy uses. Figure from Morimoto 1993.} \\label{alloc} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0002/astro-ph0002386_arXiv.txt": {
        "abstract": "We present detections of emission at 250 GHz (1.2 mm) from two high redshift QSOs from the Sloan Digital Sky Survey sample using the bolometer array at the IRAM 30m telescope. The sources are SDSSp 015048.83+004126.2 at $z = 3.7$, and SDSSp J033829.31+002156.3 at $z = 5.0$, which is the third highest redshift QSO known, and the highest redshift mm emitting source yet identified.  We also present deep radio continuum imaging of these two sources at 1.4 GHz using the Very Large Array.  The combination of cm and mm observations indicate that the 250 GHz emission is most likely thermal dust emission, with implied dust masses $\\approx 10^8$ M$_\\odot$.  We consider possible dust heating mechanisms, including UV emission from the active nucleus (AGN), and a massive starburst concurrent with the AGN, with implied star formation rates $> 10^3$ M$_\\odot$ year$^{-1}$. ",
        "introduction": "Observations  were made using MAMBO (Kreysa et al. 1999) at the IRAM 30m telescope in December 1999 and February 2000. MAMBO is a 37 element bolometer array sensitive between 190 and 315 GHz. The half-power sensitivity range is 210 to 290 GHz, and the effective central frequency for a typical dust emitting source at high redshift is 250 GHz. The beam for the feed horn of each bolometer is matched to the telescope beam of 10.6$''$, and the bolometers are arranged in an hexagonal pattern with a beam separation of 22$''$. Observations were made in standard on-off mode, with 2 Hz chopping of the secondary by 50$''$ in azimuth. The data were reduced using IRAM's NIC and MOPSI software packages (Zylka 1998). Pointing was monitored every hour, and was found to be repeatable to within 2$''$. The sky opacity was monitored every hour, with  zenith opacities between  0.23 and 0.36. Gain calibration was performed using observations of Mars, Neptune, and Uranus.  We estimate a 20$\\%$ uncertainty in absolute flux density calibration based on these observations.  The target sources were centered on the central bolometer in the array (channel 1), and the temporally correlated variations of the sky signal (sky-noise) detected in the remaining bolometers was subtracted from all the bolometer signals. The total on-target plus off-target observing time for SDSS 0338+0021 was 108 min, while that for SDSS 0150+0041 was 77 min.  The source SDSS 0338+0021 was detected with a flux density of $3.7 \\pm 0.5$mJy (Figure 1).  At $z = 5.0$, 250 GHz corresponds to an emitted frequency of 1500 GHz, or a wavelength of 200 $\\mu$m. The source SDSS 0150+0041 was detected at 250 GHz with a flux density of $2.0 \\pm 0.4$ mJy (Figure 1). At $z = 3.7$, 250 GHz corresponds to an emitted frequency of 1180 GHz, or a wavelength of 254 $\\mu$m. The quoted errors in flux density do not include the $20\\%$ uncertainty in gain calibration. The sources were not seen to vary dramatically ($\\le 30\\%$) between December 1999 and February 2000.  Assuming a dust spectrum of the type seen in the low redshift starburst galaxy Arp 220 (corresponding roughly to a modified black body spectrum of index 1.5 and temperature of 50 K), the implied 60$\\mu$m luminosity (Rowan-Robinson et al. 1997) for SDSS 0338+0021 is L$_{60} = 8.2 \\pm 1.0  \\times 10^{12}$ L$_\\odot$, while that for 0150+021 is  L$_{60} = 5.0 \\pm 1.0 \\times 10^{12}$ L$_\\odot$. Increasing the temperature to 100 K would increase the values of L$_{60}$ by a factor of about 3.5, while using the spectrum of M82 as a template would increase the values by a factor of 1.5.  The 1.4 GHz VLA observations were made in the A (30 km) and BnA (mixed 30 km and 10 km) configurations on July 8, August 14, September 30, and October 8, 1999, using a total bandwidth of 100 MHz with two orthogonal polarizations.  Each source was observed for a total of 2 hours, with short scans made over a large range in hour angle to improve Fourier spacing coverage. Standard phase and amplitude calibration were applied, as well as self-calibration using background sources in the telescope beam. The absolute flux density scale was set using observations of 3C 48.  The final images were generated using the wide field imaging and deconvolution capabilities of the AIPS task IMAGR.  The observed rms noise values on the images were within 30$\\%$ of the expected theoretical noise. The Gaussian restoring CLEAN beam was between 1.5$''$ and 2$''$ FWHM.  Neither source was detected at 1.4 GHz.  For SDSS 0338+0021 the observed value at the source position was $37 \\pm 24$ $\\mu$Jy beam$^{-1}$, and the peak in an 8$''$ box centered on the source position was 60$\\mu$Jy beam$^{-1}$ located 3.4$''$ southwest of the source position.  At $z = 5.0$, 1.4 GHz corresponds to an emitted frequency of 8.4 GHz.  For SDSS 0150+0041 the observed value at the source position was $-3.0 \\pm 19$ $\\mu$Jy beam$^{-1}$, and the peak in an 8$''$ box centered on the source position was 55$\\mu$Jy beam$^{-1}$ located 2.8$''$ north of the source position.  At $z = 3.7$, 1.4 GHz corresponds to an emitted frequency of 6.6 GHz.  We quote maximum and minimum values in an 8$''$ box in the eventuality that the dust emission is not centered on the QSO position (see discussion below). In the analysis below we adopt an upper limit of 60$\\mu$Jy beam$^{-1}$ for both sources.  Assuming a radio spectral index of --0.8, the 3$\\sigma$ upper limit to the rest frame spectral luminosity at 5 GHz of SDSS 0338+0021 is $3.3 \\times 10^{31}$ erg s$^{-1}$ Hz$^{-1}$, while that for 0150+0041 is $1.7 \\times 10^{31}$ erg s$^{-1}$ Hz$^{-1}$.  These upper limits place the sources in the radio quiet regime, using the division at $10^{33}$ erg s$^{-1}$ Hz$^{-1}$ at 5 GHz suggested by Miller et al. (1990).  Note that 90$\\%$ of optically selected high redshift QSOs are radio quiet according to this definition (Schmidt et al. 1995). ",
        "conclusions": "The fact that the cm flux densities for these sources are at least two orders of magnitude below the mm flux densities, and that the sources do not appear to be highly variable at 250 GHz, argues that the mm signal is thermal emission from warm dust.  Adopting the relation between L$_{60}$ and dust mass for hyper-luminous infrared galaxies (i.e. galaxies with L$_{60}$ $\\ge 10^{13}$ L$_\\odot$), derived by Rowan-Robinson (1999), the dust mass in SDSS 0338+0021 is $3.5 \\times 10^8$ M$_\\odot$, while that in SDSS 0150+0021 is $2 \\times 10^8$ M$_\\odot$.  This gives gas masses of $10 \\times 10^{10}$ M$_\\odot$, and $6 \\times 10^{10}$ M$_\\odot$, respectively, using the dust-to-molecular gas mass ratio of 300 suggested by Rowan-Robinson (1999).  Of course, such estimates using data at a single frequency are quite uncertain due to the lack of knowledge of the dust temperature and emissivity law, and uncertainties in the gas-to-dust ratio. We are currently searching for CO emission from these two sources to obtain a better understanding of the gas content of these QSOs.  Omont et al. (1996a) found that dust emission seems to be correlated with broad absorption line systems (BALs) in the APM QSO sample. The source SDSS 0150+0041 has been classified as a `mini-BAL' by Fan et al. (1999). The source SDSS 0338+0021 also shows strong and broad associated C IV absorption in high signal-to-noise spectra (Songaila et al. 1999). The presence of strong associated absorption is yet another indication of a rich gaseous environment in these systems.  The critical question when interpreting thermal dust emission from high redshift QSOs is whether the emission is powered by the AGN or star formation.  This question has been considered in great detail for ultra- and hyper-luminous infrared galaxies (Sanders \\& Mirabel 1996, Sanders et al. 1989, Genzel et al. 1998), and a review of this question can be found in Rowan-Robinson (1999).  Dust emitting luminous QSOs, such as SDSS 0338+0021 and SDSS 0150+0041, comprise a subset of the ultra- and hyper-luminous infrared galaxies. In QSOs, typically less than 30$\\%$ of the bolometric luminosity is emitted in the infrared. Using the blue luminosity bolometric correction factor of 16.5 derived for the PG QSO sample by Sanders et al. (1989), the FIR luminosity of SDSS 0338+0021 accounts for only 15$\\%$ of its bolometric luminosity, and the FIR emission of SDSS 0150+0041 comprises only 3$\\%$ of its bolometric luminosity. It is important to keep in mind that the FIR luminosity estimates are based on measurements at a single frequency, and hence have at least a factor two uncertainty (Adelberger \\& Steidel 2000), while the bolometric luminosity estimates require an extrapolation from the rest frame UV measurements into the blue, and hence have comparable uncertainties. Still, it is likely that only a minor fraction of the total AGN luminosity needs to be absorbed and re-emitted by dust to explain the FIR luminosity in both SDSS 0338+0021 and SDSS 0150+0041.  An important related point is that the IR emitting region has a minimum size of about 0.5 kpc for sources such as SDSS 0338+0021 and SDSS 0150+0041, as set by the far IR luminosity and assuming optically thick dust emission with a dust temperature of 50 K (Benford et al. 1999, Carilli et al. 1999).  In the Sanders et al. (1989) model for AGN-powered IR emission from QSOs, dust heating on kpc scales is facilitated by assuming that the dust is distributed in a kpc-scale warped disk, thereby allowing UV radiation from the AGN to illuminate the outer regions of the disk. Detailed models by Andreani, Franceschini, and Granato (1999) and Willott et al. (1999) show that the  dust emission spectra from 3$\\mu$m to 30$\\mu$m can be explained by such a model.  The alternative to dust heating by the AGN is to assume that there is active star formation co-eval with the AGN in these systems.  Omont et al. (1996) and Rowan-Robinson (1999) argue that star formation would be a natural, although not required, consequence of the large gas masses in these systems.  Imaging of a few high redshift dust emitting AGN shows that in some cases the dust and CO emission come from regions that are separated from the location of the optical AGN by tens of kpc (Omont et al. 1999b, Papadopoulos et al. 2000, Carilli et al. 2000).  Such a morphology argues strongly for a dust heating mechanism unrelated to UV emission from the AGN, at least in these sources. Lastly, Rowan-Robinson (1999) performed a detailed analysis of the spectral energy distributions (SEDs) of hyper-luminous infrared galaxies and concluded that, although $\\ge 50\\%$ of these systems show evidence for an AGN in the optical spectra, the SEDs between 50$\\mu$m and 1mm are best explained by dust heated by star formation. Yun et al. (2000) have extended this argument to cm wavelengths and reach a similar conclusion.  Carilli \\& Yun (2000) present a model for the expected behavior with redshift of the observed spectral index between cm and submm wavelengths for star forming galaxies, relying on the tight radio-to-far IR correlation found for nearby star forming galaxies (Condon 1992). Using the observed mm flux densities of SDSS 0150+0041 and SDSS 0338+0021, these models predict  flux densities at 1.4 GHz between 15 and 60 $\\mu$Jy for both sources. Radio images with a factor three better sensitivity are required to determine if these two sources have cm-to-mm SEDs consistent with low redshift star forming galaxies.  If the dust and radio continuum emission is powered by star formation in SDSS 0338+0021 and SDSS 0150+0041, then the star formation rates in these galaxies are high.  Using the relation between L$_{60}$ and total star formation rate given in Rowan-Robinson (1999) for a $10^8$ year starburst assuming a standard Salpeter IMF, the implied rate for SDSS 0338+0021 is 2700 M$_\\odot$ year$^{-1}$, while that for SDSS 0150+0041 is 1800 M$_\\odot$ year$^{-1}$.  Note that the high redshift QSOs from the APM sample were detected at flux levels between 3mJy and 12mJy at 240 GHz (Omont et al. 1996a), hence the required star formation rates may be even larger in some systems. The implication is that we may be witnessing the formation of a large fraction of the stars of the AGN host galaxy on a timescale $\\le 10^8$ years.  Magnification by gravitational lensing would lower the required luminosities, bringing the sources more in-line with known ultra-luminous infrared galaxies. Two of the APM QSOs detected at 240 GHz by Omont et al. (1996a) show possible evidence for gravitational lensing, while two other high redshift dust emitting QSOs, H1413+117 and APM 08279+5255, are known to be gravitationally lensed (Barvainis et al. 1994, Downes et al. 1999).  However, thus far there is no evidence for multiple imaging on arc-second scales for either SDSS 0338+0021 or SDSS 0150+0041 in optical and near IR images (Fan et al. 2000).  Imaging at sub-arcsecond resolution is required to determine if these sources are gravitationally lensed.  An important question concerning the formation of objects in the universe is: which came first, black holes or stars (Rocca-Volmerange et al. 1993)?  If the dust emission is powered by a starburst in SDSS 0338+0021 and SDSS 0150+0041, then the answer to the above question in some systems may be: both.  Co-eval starbursts and AGN at high redshift may not be surprising, since both may occur in violent galaxy mergers as predicted in models of structure formation via hierarchical clustering (Franceschini et al. 1999, Taniguchi, Ikeuchi, \\& Shioya 1999, Blain et al. 1999,  Kauffmann \\& Haehnelt 2000, Granato et al. 2000).    To properly address the interesting question of AGN versus starburst dust heating in these high redshift QSOs requires well sampled SEDs from  cm to optical wavelengths, and perhaps most importantly, imaging of the mm and cm continuum, and CO emission, with sub-arcsecond resolution.  \\vskip 0.2truein  The VLA is a facility of the National Radio Astronomy Observatory (NRAO), which is operated by Associated Universities, Inc. under a cooperative agreement with the National Science Foundation. This work was based on observations carried out with the IRAM 30 m telescope.  IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain).  This research made use of the NASA/IPAC Extragalactic Data Base (NED) which is operated by the Jet propulsion Lab, Caltech, under contract with NASA.  CC acknowledges support from the Alexander von Humboldt Society.  DPS acknowledges support from National Science Foundation Grant AST99-00703.  XF and MAS acknowledge support from the Research Corporation, NSF grant AST96-16901, and an Advisory Council Scholarship."
    },
    "0002/astro-ph0002179_arXiv.txt": {
        "abstract": " ",
        "introduction": "As yet, little is known about the structure of the inner regions of powerful radio galaxies and the impact of the activity on circumnuclear regions. The high resolution and sensitivity afforded by imaging observations using the Hubble Space Telescope (HST) has revealed a wealth of complex structures in such objects, with dust lanes, jets and large regions of scattered emission from hidden nuclei. One of the key targets in studies which aim to elucidate the interplay between active nuclei and their host galaxies, and between galactic activity of different types, is the archetypal powerful radio galaxy Cygnus~A (see Carilli \\& Barthel 1996 for a review).  HST infrared imaging observations of Cygnus~A by Tadhunter \\etal\\ (1999) revealed an edge-brightened bi-conical structure centred on the nuclear point source, which is strikingly similar to those observed around young stellar objects. The edge-brightening of this structure provides evidence that the bicone is defined as much by outflows in the nuclear regions as by the polar diagram of the illuminating quasar radiation field. The HST observations also show an unresolved nuclear source at 2.0 and 2.25$\\mu$m. However, from the imaging observations alone it is unclear whether this unresolved source represents the highly extinguished quasar nucleus seen directly through the obscuring torus, or emission from a less-highly-extinguished extended region around the nucleus.  Near-IR polarization observations have the potential to remove the uncertainties surrounding the nature of the unresolved  nuclear sources and, in addition, to provide further important information about the obscuration and anisotropy in the near-nuclear regions of powerful radio galaxies. Previous ground-based polarimetric observations of Cygnus~A by Packham \\etal\\ (1998) demonstrate that the nuclear regions are highly polarized in the K-band, with a measured polarization of $P_k \\sim 4$\\% for a 1 arcsecond diameter aperture centred on the compact IR nucleus. However, the resolution of the ground-based observations is insufficient to resolve the structures and determine the polarization mechanism unambiguously. In this letter we present new diffraction-limited infrared imaging polarimetry observations of Cygnus~A made with the Near Infrared Imaging Camera and Multi-Object Spectrometer (NICMOS: MacKenty \\etal\\ 1997) on the HST. These observations resolve the polarized structures, and raise new questions about the nature of the anisotropy in the near-nuclear regions of this key source. ",
        "conclusions": "\\subsection{The nature of the unresolved core source}  A major motivation for the HST observations was to investigate the nature of the compact core source and the cause of the relatively large polarization measured in the core by Packham \\etal\\ (1998). The explanation favoured by Packham \\etal\\ is that the compact core source represents transmitted quasar light, while the high polarization is due to dichroic absorption by aligned dust grains in the central obscuring torus. Because the dichroic mechanism is relatively inefficient, a high polarization implies a large extinction: from observations of Galactic stars it is known that {\\it at least} 55 magnitudes of visual extinction is required to produce a K-band polarization of 28\\% for optimum grain alignment (Jones 1989). More typically, the correlation between K-band polarization and extinction deduced for Galactic stars by Jones (1989) implies that an extinction of $A_v \\sim 350$ magnitudes would be required for $P_k = 28$\\%. For comparison, an upper limit on the K-band extinction in Cygnus A, estimated by comparing the 2.25$\\mu$m core flux  with mid-IR and X-ray fluxes, is $A_v < 94$ magnitudes (see Tadhunter \\etal\\ 1999 for details).  Thus, the dichroic mechanism is only feasible if the efficiency of the mechanism in Cygnus A is greater than it is along most lines of sight in our Galaxy. Such enhanced efficiency cannot be entirely ruled out, given that the  Galactic dichroic polarization involves a randomly oriented magnetic field component, whereas the magnetic fields in the central obscuring regions of AGN may be more coherent. In this context it is notable that near- and mid-IR polarization measurements of the central regions of the nearby Seyfert galaxy NGC1068 provide evidence for a greater dichroic efficiency than predicted by the Jones (1989) correlation, with $P_k = 5$\\% produced by $A_v \\sim$20 -- 40 magnitudes (Lumsden et al. 1999). However, even the greater dichroic efficiency deduced for NGC1068 would not be sufficient to produce the high polarization measured in the core of Cygnus A if $A_v < 94$ magnitudes.  The efficiency problem might be resolved if the extinction to the core source in the K-band is higher than the $A_v < 94$ estimated on the basis of comparisons of the K-band flux with the mid-IR and X-ray fluxes. Indeed, substantially higher extinctions have been deduced for Cygnus A, both from modelling the X-ray spectrum of the core ($A_v = 170\\pm30$: Ueno \\etal\\ 1994) and from comparisons between hard-X-ray continuum, [OIII] emission line and mid-infrared continuum fluxes ($A_v = 143\\pm35$: Ward 1996, Simpson 1995). If such high extinctions also apply to the quasar nucleus in the K-band, the low efficiency of the dichroic mechanism would be less of a problem. However, for any reasonable quasar SED, the contribution of such a highly obscured quasar nucleus to the flux and the polarization of the {\\it detected} 2.0$\\mu$m core source would be negligible (i.e. we would not expect to detect the quasar nucleus directly in the K-band). Thus, it is more likely that the relatively low extinction deduced from the K-band flux measurements reflects contamination of the K-band core by emission from a less-highly-obscured region, which is close enough to the central AGN to remain  unresolved at the resolution of our HST observations. Although it has been proposed that the contaminating radiation in the K-band may include hot dust emission and/or line emission from quasar-illuminated regions close to the nucleus (e.g. Stockton \\& Ridgway 1996), such emission would have a low intrinsic polarization, and a large dichroic efficiency would still be required in order to produce the polarization of this component by dichroism.  The most plausible alternative to dichroic extinction is that the K-band core source represents scattered- rather than transmitted quasar light. In this case, the polarization is a consequence of scattering in an unresolved region close to the illuminating quasar; we do not detect the quasar nucleus directly in the K-band; and previous extinction estimates based on the K-band fluxes substantially underestimate the true nuclear extinction. Note that the presence of such a scattered component would resolve the discrepancy between the extinction estimates based on K-band flux measurements, and those based on fluxes measured at other wavelengths.      Finally, we must also consider the possibility that the core polarization is due to synchrotron radiation associated with the pc-scale jet visible in VLBI radio images (Krichbaum \\etal\\ 1996). Although the {\\it integrated} polarization of the radio core is small even at high radio frequencies ($P_{22GHz} < 5$\\%: Dreher 1979), we cannot entirely rule out the possibility that we are observing a highly polarized sub-component of the jet which suffers a relatively low extinction, or alternatively that the radio core source as a whole suffers large Faraday depolarization at radio wavelengths, and would appear more highly polarized at infrared wavelengths. Polarization observations at sub-mm wavelengths will be required to investigate this latter possibility.   \\subsection{The extended polarization structures}  An intriguing feature of our HST observations is the high degree of polarization measured along, and only along, the NW-SE limb of the bicone. The orientation of the polarization measured along the limb is consistent with the scattering of light from a compact illuminating source in the nucleus, while the high degree of polarization is consistent with the edge-brightened bi-cone geometry of Tadhunter \\etal\\ (1999), in the sense that the scattering angle for the edge-brightened region will be close to the optimal 90$^\\circ$ required for maximal polarization. However, the fact that the polarization is measured along only one limb is difficult to reconcile with the simplest bicone model in which the illuminating IR radiation field is azimuthally isotropic, and the scattering medium is uniformly distributed around the walls of the funnels hollowed out by the circum-nuclear outflows. In this simplest model both limbs would be highly polarized in the direction perpendicular to the radius vector of the source, and this is clearly inconsistent with the observations.  Our observation require that one or more of the assumptions implicit in the simple model must be relaxed. In general terms this requires either invoking specific matter distributions within the cone and/or an anisotropic illumination pattern of the central source itself.  Perhaps the simplest way of reconciling the polarization characteristics with the bi-cone geometry is to adjust the relative importance of scattering and intrinsic emmission with azimuth around the cone, so that one limb of the cone is dominated by scattered radiation, while the other is predominatly intrinsic radiation. Since the band-pass of the NICMOS polarizers contains the Paschen alpha line, this provides an obvious potential source of the diluting radiation for the unpolarized regions. Unfortunately there is no obvious reason why such an asymmetry should exist. Furthermore, direct images with the F222M filter show that there is no radical change in the relative brightness of the two limbs of the eastern cone between 2.0$\\mu$m and 2.25$\\mu$m. This is an argument against  the Paschen alpha model for the intrinsic emission, since the F222M filter admits no emission lines as strong as Paschen alpha.   An alternative possibility is that the NW-SE limb is brighter because the near-IR radiation field of the illuminating AGN is more intense in that direction (i.e. the illuminating radiation field is azimuthally anisotropic within the cones). In this case, the clear difference in structure between the optical and near-IR reflection nebulae suggests that the sources of illumination at the two wavelengths are different: while the source of the shorter wavelength continuum must produce a radiation field which is azimuthally isotropic within cones defined by the obscuring torus, the source of illumination at the longer wavelengths is required to display considerable degree of azimuthal anisotropy. The near-IR continuum source must also have a relatively red spectrum, in order to avoid producing similar structures at optical and infrared wavelengths. The near-IR anisotropy might arise in the following ways. \\begin{enumerate} \\item {\\bf Beamed radiation from the inner radio jet.} The near-IR continuum is emitted by a component of the inner synchrotron jet which has a direction of bulk relativistic motion siginificantly displaced from the axis of the large-scale radio jet, such that the radiation is beamed towards the NW-SE limb of the bicone. However, given the remarkable degree of collimation, and the lack of bending, observed in the Cygnus~A jet on scales betweem 1pc and 100kpc, a major difficulty with this model is that the the jet would have to bend through a large angle on a scale smaller than $\\sim$1pc (the resolution of the VLBI maps), whilst retaining the rotational symmetry in the jet structure about the nucleus. A further requirement of this model is that, if the inner jet is precessing, the precession timescale must be greater than the light travel time ($\\sim$5000 years) across the bicone structure. \\item {\\bf Anisotropic hot dust emission.} The near-IR continuum is emitted by hot dust in the inner regions of the galaxy, with a larger projected area of the emitting region visible from one limb of the bicone than from the other. For example, if the near-IR radiation is emitted by dust in a warped disk close to the central AGN --- perhaps an outer part of the accretion disk --- the warp could be oriented such that the NW-SE limb has an almost face-on view of the emitting region, whereas the the NE-SW limb has a more oblique view. There is already direct observational evidence for a warped outer accretion disk in at least one active galaxy (Miyoshi \\etal\\ 1995). Given that this mechanism would produce a relatively mild, broad-beam anisotropy, an optically thick torus on scales larger than the hot dust emitting region would still be required in order to produce the sharp-edges to the illuminated bicone structure. \\end{enumerate} Note that regardless of how any anisotropy in the IR radiation field might be produced, such anisotropy would not by itself explain the nature of the unpolarized emission along the SW-NE limb of the bicone, and the lack of variation in the brightness ratio of the two limbs of the eastern cone between 2.0$\\mu$m and 2.25$\\mu$m. It may also be difficult to reconcile the anisotropic illumination model with the polarization properties of the unresolved core source: if the core polarization is due to scattering, the orientation of the core polarization vector implies that a substantial flux of illuminating photons must escape at large angles to the NW-SE limb of the bicone.   We expect future spectropolarimetry observations to resolve the uncertainties concerning the origin of the near-IR polarization structures in Cygnus~A. For example, in the case of the anisotropic illumination mechanisms considered above, the anisotropy is in the {\\it continuum} flux rather than the broad lines associated with the AGN. Thus,  if this model is correct, the broad lines will be relatively weak or absent in the polarized spectrum of the extended structures. In contrast, for a non-uniform distribution of scattering material but isotropic illumination within the cones, the broad lines and continuum will be scattered equally, and the equivalent widths of the broad lines in the polarized spectrum should fall within the the range measured for steep spectrum radio quasars."
    },
    "0002/hep-th0002070_arXiv.txt": {
        "abstract": "{We propose a picture, within the pre-big-bang approach, in which the universe emerges from  a bath of {\\em plane} gravitational and dilatonic waves. The waves interact gravitationally breaking the exact plane symmetry and lead generically to gravitational collapse resulting in a singularity with the Kasner-like structure. The analytic relations between the Kasner exponents and the initial data are explicitly evaluated and it is shown that pre-big-bang inflation may occur within a dense set of initial data. Finally, we argue that plane waves carry zero gravitational entropy and thus are, from a thermodynamical point of view, good candidates for the universe to emerge from.} ",
        "introduction": "The low energy effective equations of string theory provide cosmological solutions which might be applicable just below the string scale in the very early universe. In the pre-big-bang (PBB) scenario, suggested naturally by the spirit and the symmetries of Superstring theory, the universe starts in a low curvature, low coupling regime and then enters a stage of dilaton driven kinetic inflation \\cite{gasp}.  To address one of the main problems of cosmology, namely the problem of the initial conditions, this interesting picture has been  developed further in \\cite{buon}, where the authors suggest that  the initial state of the universe could have  consisted of a bath of gravitational and dilatonic waves, some of which would have  collapsed leading  to the birth of a baby inflationary universe. These PBB bubble universes would give rise, finally, after a yet-to-be clarified graceful exit mechanism, to the observed Friedman-Robertson-Walker (FRW) world.  The main purpose of this paper is to develop a modified {\\em realization} of the PBB bubble picture of Buonanno,  Damour and Veneziano \\cite{buon}, in which the spherically symmetric collapse leading to inflationary PBB solutions is substituted by the interaction of strictly plane waves. This modification affects only the initial state of the universe, while near the (spacelike) caustic singularity the model shows similar behaviour to that discussed in \\cite{buon}, leading  to Kasner-like structure. Representing exact solutions to the classical string equations of motion to all orders in the inverse string tension \\cite{horstief}, it looks rather natural to modify the PBB picture by incorporating plane waves into the postulate of ``asymptotic past triviality\" \\cite{buon}. Moreover, this picture is  attractive not only due to the exactness of plane wave backgrounds for string propagation, but  most importantly, because of mutual ``fatal attraction\" exercised by the plane waves which leads to an inevitable  gravitational collapse independently of their strength, unlike in the spherical picture.  In the scenario we propose, one starts with a model universe in a low coupling, low curvature regime   with plane gravitational and matter waves which eventually will gravitationally interact \\footnote{In the gravitational sector we limit ourselves to the discussion of constantly polarized waves (diagonal metrics) on the grounds that in string theory the collision of gravitational waves with variable polarization may be generically mapped via T-duality into another problem where the variable polarization of the incoming waves is transformed into an non-vanishing value of the $B_{\\mu\\nu}$ field \\cite{flvm}.}. Colliding plane wave space-times have been investigated in detail in general relativity (see \\cite{grif} and references therein). A generic feature of the interaction of two plane waves is the formation of a strong space-like curvature singularity in the future \\cite{tip}. In the context of the PBB scenario, this singularity can be re-interpreted as  an ordinary cosmological singularity.  The approach to the singularity from the past occurs through a Kasner-like behaviour, and this, in  turn, can be analytically related to the initial data. This will provide the framework for a future quantitative study of the important problem of fine tuning of initial conditions \\cite{turn,jrd,kl,fvm} leading to {\\em inflationary} behaviour as $t\\rightarrow 0^{-}$. This problem is well defined in our picture, unlike  in the spherically symmetric case where a similar analysis does not seem to be possible. In the case at hand, the problem is a direct generalisation of the problem of determining Kasner exponents in scattering of pure plane gravitational waves with constant polarization \\cite{yurt}. As far as the technical part of this paper is concerned, some already  well-known results from general relativity  and mathematical cosmology will be used and re-interpreted in a new light.  The whole question of naturalness of initial conditions is far from being settled and is under current discussion in the literature. The authors of \\cite{kl} have found that the solution of the flatness problem requires the introduction of two huge dimensionless parameters. On the other hand in \\cite{gas} it was argued that this might not constitute such a fine tuning if the initial scale is taken to be the whole horizon and not just the Planck scale. From a different point of view the genericity of PBB cosmology has been addressed in \\cite{jrd,gr}, showing that plane waves might be considered a generic initial state of PBB cosmology. In this paper, however, we will not enter into the discussion of whether fine tuning is needed in order to solve the flatness problem, but will rather stick to the possibility of inflation, leaving the resolution of these other questions to future work.  The paper is organized as follows. In Sec. 2 we review briefly the initial value problem for the collision of dilatonic and gravitational waves. In particular we  provide a closed expression relating the Kasner exponents which characterize the asymptotic geometry near the caustic singularity with the initial conditions for the metric functions and the dilaton. In Sec. 3 we  apply these results to investigate the range of initial conditions leading to PBB inflation, and whether these conditions are naturally met in the collision of plane gravitational waves. In Sec. 4 we study two particular geometries in the interaction region, the Nappi-Witten solution \\cite{nappi}, and a family of Kantowski-Sachs metrics studied in \\cite{fvm}. Finally, in Sec. 5 we will use thermodynamical considerations to argue that plane waves are good candidates to represent the primordial PBB universe, summarize our conclusions and indicate possible future directions of research. ",
        "conclusions": "In this  paper we have proposed a picture where the PBB inflation is realized starting from a trivial asymptotic state. The idea is to start with  strictly plane waves moving in different directions which interact  gravitationally at some stage to produce a space-time singularity. The structure of the solutions close to the singularity is of the Kasner type, and we were able to relate analytically the initial data on the wave fronts to the Kasner exponents in the string frame. Although maybe the collapsing regions studied here are not as generic as for example those studied in \\cite{buon}, our picture has two basic advantages: \\begin{itemize} \\item[1.] The initial background is exact from the point of view of string propagation. \\item[2.] One can completely determine the structure of the singularity in terms of the initial data provided by the incoming waves. \\end{itemize} These two elements make of this scenario, at least, a solid test bench and a sort of theoretical laboratory for the PBB ideas.  We have seen that there exists a dense set of initial data leading to inflationary behaviour in the PBB phase and, therefore, there is a good chance for an inflationary universe to emerge during this phase.  Since the Kasner exponents actually carry a space dependence, i.e. they depend in general on one coordinate $z$, different regions with different Kasner exponents experience different types of inflation. Therefore, although we start with the collision of two plane waves we obtain near the singularity a rich structure with the formation of different multiple PBB inflationary bubbles.  Although we have extracted these conclusions from a general analysis of the collision of two gravi-dilatonic waves, we have also studied two concrete examples of initial conditions leading to different geometries in the interaction region. The first one leads to a Nappi-Witten solution in region IV and it has the obvious interest of providing an exact string background also in this region. The second example corresponds to a spherically symmetric interaction region that may describe gravitational collapse as the result of the collision.  Another motivation to look at plane wave space-times as a probable initial state for the PBB universe may come from thermodynamical considerations. One expects the universe to start in the lowest possible state of gravitational entropy. Let us suppose, arguing in the spirit of Penrose's hypothesis \\cite{penrose}, that we relate the gravitational entropy to some homogeneous function constructed from all possible curvature invariants and that we normalize this function to be vanishing when the invariants vanish. With such a function at hand we formulate a kind of ``Generalised Third Law of Thermodynamics'' and assign zero gravitational entropy to those spacetimes for which {\\em all} curvature invariants vanish. Interestingly, the FRW models do not fall into the class of zero entropy models according to our definition  unlike in the Penrose's case. To justify this, we argue  that the lowest entropy state, apart from being  the simplest, must be an exact string background. The space-times with all vanishing curvature invariants ({\\it plane waves}) certainly do so, while the FRW universe does not.  The additional support to consider the gravitational entropy content of the plane wave geometry to be vanishing comes from yet a different, though not unrelated argument. One would usually tend to relate the gravitational entropy with the phenomenon of quantum particle creation. It is commonly accepted that quantum particle creation indicates whether a system is endowed with nontrivial gravitational entropy. Plane waves, due to their symmetry and to the fact that all the curvature invariants vanish, do not polarize the vacuum, so quantum particles are not created in the vicinity of plane waves \\cite{des}. This is consistent then with the hypothesis of assigning zero gravitation entropy to the plane wave. Moreover, it looks as if time is not a player in the plane wave regime. Due to the absence of the global Cauchy surface, one may consider such a pure plane wave geometry as ``timeless''. Therefore, until two such  waves interact, no notion of time as defined by entropy change is appreciated. What happens further is beyond the scope of this paper.  One of the most interesting issues, untouched here, is the problem of whether the gravitational wave collision problem can be {\\it globally} defined in terms of an exact string background. Provided one starts with initial states that are exact string backgrounds, what are the conditions for the data to evolve into the interaction region without breaking conformal invariance? We know that one such solution exists in this region, namely the Nappi-Witten solution that we studied in Sec. 4.1. Therefore, since we start with plane gravitational waves (which are exact string backgrounds) and we can make the transition over the null boundaries as much differentiable as we like, the question remains of whether this implies conformal invariance in the interaction region. On purely physical grounds one would be tempted to say that exact string backgrounds in the regions II and III of Fig. 1 will smoothly (i.e. $C^{\\infty}$) extend to an exact background in region IV, at least if the full string equations share the uniqueness properties of the Einstein equations.  An interesting issue to address would be to refine the picture provided here to account for more realistic situations in which the primordial gravitational waves are not plane but localized. It has been argued that for ``almost'' plane gravitational waves singularities also occur as the result of their collision \\cite{yuapw}. In the case of ``graviton beams'', there is also mutual focusing and maybe production of singularities \\cite{vgb} that could serve as seeds for PBB bubbles.  Finally, there has been some work on the collision of plane waves at Planckian energies \\cite{planck} leading to black hole nucleation through tunneling. This kind of scenarios might be a way to find a semi-classical approximation of the formation of the singularity that could be applied to the graceful exit problem in PBB cosmology. This, and the thermodynamical ideas we have outlined above, we hope to be able to discuss in the future."
    },
    "0002/astro-ph0002447_arXiv.txt": {
        "abstract": "We present near-IR photometry and imaging observations of a small sample of sources identified in the BeppoSAX 5--10 keV survey (HELLAS) which resolves $\\sim 20-30$\\% of the X-ray background at these energies. The near-IR data are combined with optical spectra and photometry. Only 40\\% of the sources in our sample have the blue, power law continuum typical of color--selected QSOs. The remaining 60\\% are dominated by a galactic component which, on the basis of the continuum colors and shape, have ages ranging from 10$^9$ to 10$^{10}$ years. The images show that the blue QSOs are pointlike at our angular resolution, while all the other sources are extended, consistent with their spectral appearance and low redshift. Since down to R~=~20 only about two thirds of the HELLAS sources have a counterpart, the preliminary HELLAS census comprises in roughly equal parts: i) blue QSOs (mostly at high redshifts); ii) optically dim, galaxy--dominated active nuclei (mostly at modest redshifts); and iii) empty fields (possibly highly absorbed QSOs at high redshifts). ",
        "introduction": "X-ray background (XRB) synthesis models ascribe most of the high energy flux to radio quiet, absorbed Active Galactic Nuclei (AGNs) at intermediate and high redshifts (e.g. Comastri et al. 1995, Gilli et al. 1999, Pompilio et al. 2000). Observationally the available information on AGNs at these redshifts refers mostly to unabsorbed nuclei, since the current samples of radio quiet AGNs have been selected mostly with color techniques, or in the soft X~rays. Only recently have selection criteria less sensitive to absorption been used. Examples are the radio quiet red QSOs (Kim \\& Elvis 1999), analogous to the ones already found in radio loud samples (eg. Webster et al. 1995), and the spectroscopic identifications in the ELAIS field (Rowan--Robinson et al. 1999). Yet, most of our knowledge about absorbed, radio quiet AGNs is limited to low redshifts and low luminosities, where spectroscopic surveys of bright galaxies have been performed.  The High Energy LLarge Area Survey [HELLAS, Comastri et al. 2000, Fiore et al. 2000 (paper II)] aims at providing a useful sample of hard X-ray selected (5--10 keV), optically identified AGNs while waiting for the Chandra and XMM results. The survey instrument is the BeppoSAX MECS. The sky coverage is 1--50 square degrees at $F_{5-10~keV} =$~5--30~$\\times 10^{-14} \\rm erg~cm^{-2} s^{-1}$, respectively, and is 84 deg$^2$ at fluxes higher than $9\\times 10^{-13} \\rm erg~cm^{-2} s^{-1}$. The cataloged sources amount to 147, and at the fainter limit the source density is $16.9\\pm6.4$~deg$^{-2}$, implying that about 20--30\\% of the XRB at these energies has been resolved. A program of optical identification is underway, which includes optical/near-IR broad band photometry and near-IR imaging, beside optical spectroscopy of all candidates down to R = 20. So far, 63 optical counterparts have been identified, in about two thirds of the examined errorboxes. About half of the spectra are typical of QSOs, with a blue continuum and broad lines, about half are of intermediate type (1.8--1.9), generally with red continua, and a few of them contain only narrow lines [Fiore et al. 1999 (paper I), La Franca et al. 2000 (paper III)].  In this paper we present and discuss preliminary results of the near-IR photometry and imaging observations of the spectroscopically identified counterparts. Combined with the optical information presented in papers I and III these data give a broad-band view of the properties of the HELLAS sources and allow a preliminary census of the hard XRB contributors.  \\begin{table}[!] \\caption[]{Observing log.}\\label{tab_obs} \\begin{tabular}{cccccc} \\hline \\hline \\# &\\hskip-0.2truecm SAXJ & Date &\\hskip-0.2truecm Seeing\\hskip-0.4truecm &\\multicolumn{2}{c}{Sensit.(T$_{int}$)$^a$} \\\\ &  &  &\\hskip-0.2truecm($''$)\\hskip-0.4truecm &J & Ks \\\\ \\hline 1  & \\hskip-0.2truecm 1118.8+4026A & 19/02/99 &\\hskip-0.2truecm 1.0\\hskip-0.4truecm & 21.2(4.0) & 19.1(4.0) \\\\  % 2  & \\hskip-0.2truecm 2044.6-1028 & 22/07/99 &\\hskip-0.2truecm 0.4\\hskip-0.4truecm & 21.8(1.6) & 21.0(1.6) \\\\  % 3  & \\hskip-0.2truecm 1117.8+4018 & 19/02/99 &\\hskip-0.2truecm 0.9\\hskip-0.4truecm & 21.3(4.0) & 18.9(2.0) \\\\  % 4  & \\hskip-0.2truecm 1118.8+4026B & 19/02/99 &\\hskip-0.2truecm 1.0\\hskip-0.4truecm & 21.2(4.0) & 19.1(4.0) \\\\  % 5  & \\hskip-0.2truecm 0045.7-2515 & 18/12/98 &\\hskip-0.2truecm 1.6\\hskip-0.4truecm & 20.9(6.3) & 19.0(8.5) \\\\  % 6  & \\hskip-0.2truecm 1528.8+1939 & 02/03/99 &\\hskip-0.2truecm 1.0\\hskip-0.4truecm & --      & 19.3(5.8) \\\\  % 7  & \\hskip-0.2truecm 1519.5+6535 & 02/03/99 &\\hskip-0.2truecm 0.8\\hskip-0.4truecm & --      & 19.4(4.0) \\\\  % 8  & \\hskip-0.2truecm 1353.9+1820 & 19/02/99 &\\hskip-0.2truecm 1.0\\hskip-0.4truecm & 20.8(2.0) & 18.9(3.0) \\\\  % 9  & \\hskip-0.2truecm 1218.9+2958 & 19/02/99 &\\hskip-0.2truecm 0.9\\hskip-0.4truecm & 20.9(2.0) & 18.9(2.0) \\\\  % 10 & \\hskip-0.2truecm 1118.2+4028 & 19/02/99 &\\hskip-0.2truecm 1.0\\hskip-0.4truecm & 20.8(2.0) & 18.8(2.0) \\\\  % \\hline \\end{tabular} \\vskip0.1truecm $^a$ 3$\\sigma$ limiting magnitude (within an aperture twice the seeing) and integration time in minutes. \\end{table}  \\begin{table*}[!] \\caption[]{Results.}\\label{tab_obs} \\begin{tabular}{lccccccccccccc} \\hline \\hline \\# & F$_X^a$ & logN$_H^b$ & L$_X^c$ & type$^d$ & z & J$^e$ & Ks$^e$ & \\hskip-0.2truecm ext.$^f$ & \\hskip-0.2truecm FWHM$^g$ & B--R & B--Ks & R--Ks \\\\ \\hline 1  & 1.2 & $<22.5$ & 45.1 &  B  & 1.13 & 17.98$\\pm$0.10 & 16.69$\\pm$0.15 & \\hskip-0.2truecm N & \\hskip-0.2truecm $1''.3$(1.3)   & 0.42$\\pm$0.28 & 1.69$\\pm$0.29 & 1.27$\\pm$0.29  \\\\  % 2  & 2.0 & $>22.7$ & 46.6 &  B  & 2.76 & 16.25$\\pm$0.05 & 14.91$\\pm$0.05 & \\hskip-0.2truecm N & \\hskip-0.2truecm $0''.4$(1.0)    & 0.43$\\pm$0.28 & 2.88$\\pm$0.25 & 2.45$\\pm$0.25  \\\\  % 3  & 1.3 & $22.7^{+0.5}_{-0.5}$& 45.4 &  B  & 1.27 & 19.38$\\pm$0.10 & 17.74$\\pm$0.15 & \\hskip-0.2truecm N & \\hskip-0.2truecm $1''.0$(1.1)    & 0.55$\\pm$0.28 & 2.96$\\pm$0.29 & 2.41$\\pm$0.29   \\\\  % 4  & --$^h$ & $<22.4$ & --$^h$ &  B  & 0.89 & 19.21$\\pm$0.10 & 17.44$\\pm$0.15 & \\hskip-0.2truecm N & \\hskip-0.2truecm $1''.4$(1.4)    & 0.75$\\pm$0.28 & 3.23$\\pm$0.29 & 2.48$\\pm$0.29    \\\\  % 5  & 3.5 & $22.6^{+0.5}_{-0.5}$ & 43.3 & 1.9 & 0.11 & 16.38$\\pm$0.05 & 15.27$\\pm$0.07 & \\hskip-0.2truecm Y & \\hskip-0.2truecm $5''.1$(3.2)    & 0.80$\\pm$0.28 & 2.83$\\pm$0.25 & 2.03$\\pm$0.25    \\\\  % 6  & 1.4 & $<22.6$ & 43.3 &  L  & 0.18 & --             & 15.94$\\pm$0.07 & \\hskip-0.2truecm Y & \\hskip-0.2truecm $8''.9$(8.9)      & 1.22$\\pm$0.28 & 4.30$\\pm$0.25 & 3.08$\\pm$0.25    \\\\  % 7  &  11. & $23.2^{+0.1}_{-0.1}$ & 43.2 & 1.9 & 0.04 & --             & 11.38$\\pm$0.07 & \\hskip-0.2truecm Y & \\hskip-0.2truecm $7''.7$(7.7)  & 1.59$\\pm$0.04 & 4.60$\\pm$0.08 & 3.01$\\pm$0.07     \\\\  % 8  &  8.5 & $22.8^{+0.3}_{-1.2}$& 44.2  &  R  & 0.22 & 15.26$\\pm$0.05 & 13.90$\\pm$0.07 & \\hskip-0.2truecm Y & \\hskip-0.2truecm $3''.6$(3.6)     & 2.18$\\pm$0.04 & 5.53$\\pm$0.08 & 3.35$\\pm$0.07     \\\\  % 9  &  2.4 & $23.1^{+0.4}_{-0.4}$& 43.6 & 1.9 & 0.18 & 16.51$\\pm$0.05 & 15.14$\\pm$0.07 & \\hskip-0.2truecm Y & \\hskip-0.2truecm $2''.8$(3.1)    & 2.35$\\pm$0.28 & 5.76$\\pm$0.26 & 3.41$\\pm$0.26     \\\\  % 10  & 0.9 & $<22.1$ & 43.9 &  R  & 0.39 & 16.48$\\pm$0.05 & 14.97$\\pm$0.07 & \\hskip-0.2truecm Y & \\hskip-0.2truecm $2''.9$(2.9)   & 2.66$\\pm$0.28 & 5.83$\\pm$0.26 & 3.17$\\pm$0.26     \\\\  % \\hline \\end{tabular} \\vskip0.1truecm $^a$ 5-10 keV flux in units of $\\rm 10^{-13}~ erg~ cm^{-2} ~ s^{-1}$; $^b$ log of intrinsic absorbing column density in units of cm$^{-2}$, assuming an intrinsic photon index $\\Gamma =1.8$ (see paper I) and the absorber at the redshift of the source; $^c$ log of the intrinsic (deabsorbed) 5-10 keV luminosity in units of erg s$^{-1}$, assuming H$_0$=50 and q$_0$=0 (see paper I); $^d$ optical spectral type: 1.9~-- narrow line spectrum with broad component of H$\\alpha$, L~-- LINER, B~-- broad line QSO with blue continuum, R~-- broad line QSO with red continuum; $^e$ Photometry extracted with a circular aperture down to the noise and corrected for Galactic extinction. $^f$ Y if resolved in the near-IR images; $^g$ full-width at half-maximum along the major axis in arcseconds (quantities in brackets are in seeing units); $^h$ in the same errorbox of \\#1. \\end{table*} ",
        "conclusions": "We presented new near-IR (J and Ks band) observations of a sample of 10 objects selected in the hard X-rays (5--10 keV). These sources were discovered in a large survey (HELLAS) performed by the BeppoSAX satellite, which resolves $\\sim 20-30$\\% of the hard X-ray background. The sample includes 4 blue broad line QSOs and 6 AGNs with redder continua whose optical emission line spectra range from broad line objects (red QSOs), to intermediate type 1.9 AGNs, to LINER.  The B--K color ranges from the standard value of $\\sim 2$ (typical of U--B color selected QSOs) up to $\\sim$6, similar to the color of red QSOs found in radio surveys.  The red AGNs show extended near-IR images. Model fitting of the photometry and spectral data shows that all and only the red AGNs are {\\it dominated by the emission of the host galaxy} (with an age of 10$^9$--10$^{10}$ yr). Red AGNs amount to about a third of the total HELLAS sources, and their fraction is expected to increase significantly at fainter X-ray fluxes, where most of the hard X-ray background is produced. Therefore, our result suggests that a significant fraction of the counterparts of the sources making the hard X-ray background appear as ``normal'' galaxies at optical and near-IR wavelengths. Chandra and XMM are expected to discover a large number of this class of objects."
    },
    "0002/astro-ph0002392_arXiv.txt": {
        "abstract": "A phenomenological model for chemical enrichment in early-type galaxies is presented, in which the process of star formation is reduced to a set of four parameters: star formation efficiency ($\\ceff$), fraction of ejected gas in outflows ($\\bout$), formation redshift ($z_F$) and infall timescale ($\\tau_f$). Out of these four parameters, only variations of $\\bout$ or $\\ceff$ can account for the color-magnitude relation. A range of outflows results in a metallicity sequence, whereas a range of star formation efficiencies will yield a mixed age + metallicity sequence. The age-metallicity degeneracy complicates the issue of determining which mechanism contributes the most (i.e. outflows versus efficiency). However, the determination of the slope of the correlation between mass-to-light ratio and mass in clusters at moderate or high redshift will allow us to disentangle age and metallicity. ",
        "introduction": "One of the long-standing problems in astrophysics is the process of star formation in galaxies. The standard scenario assumes stars to form from gas that falls in the potential wells of dark matter halos. Subsequent interacting or merging stages among galaxies might trigger additional bursts of star formation. The complex nature of star formation makes this problem rather an untractable one from an analytical point of view, so that the best approach towards understanding the distribution of stellar populations in galaxies requires a heavy use of rough approximations and all-too-often dangerous generalizations. It is the purpose of current phenomenological models describing the formation and evolution of the stellar component in galaxies to reveal the mechanisms which describe the wide range of galaxy colors and luminosities as well as their connection to morphology. The current status of the determination of the ages of the stellar populations in galaxies is rather controversial due to the degeneracy between age and metallicity (Worthey 1994). Observations of early-type galaxies by two different groups using similar techniques targeting narrow spectral indices to infer a luminosity-weighted age give contradictory results. While Trager et al. (2000) find a large age spread in the sample of field and group early-type systems of Gonz\\'alez (1993), Kuntschner (2000) reports a large metallicity spread in Fornax cluster ellipticals. So far, any observational measurement of age is plagued by many degeneracies which render a direct estimate uncertain. An alternative approach modelling the formation and chemical enrichment of the stellar component of galaxies is needed in order to reveal the actual scenario of galaxy formation.  \\begin{figure} \\epsfxsize=2.6in \\begin{center} \\leavevmode \\epsffile{ferrerasi_f1.eps} \\end{center} \\vskip -0.3truein \\caption{Gas ({\\sl top}) and metallicity ({\\sl bottom}) evolution of two models: one with a series of short bursts ({\\sl thin line}) and another with a single, more extended burst. The inset shows the evolution of $U-V$ color in the rest frame. }\\label{f1} \\end{figure} ",
        "conclusions": ""
    },
    "0002/astro-ph0002501_arXiv.txt": {
        "abstract": "Study of radio supernovae (RSNe) over the past 20 years includes two dozen detected objects and more than 100 upper limits. From this work we are able to identify classes of radio properties, demonstrate conformance to and deviations from existing models, estimate the density and structure of the circumstellar material and, by inference, the evolution of the presupernova stellar wind, and reveal the last stages of stellar evolution before explosion.  It is also possible to detect ionized hydrogen along the line of sight, to demonstrate binary properties of the stellar system, and to show clumpiness of the circumstellar material. More speculatively, it may be possible to provide distance estimates to radio supernovae.  The interesting and unusual radio supernova SN 1998bw, which is thought to be related to the $\\gamma$-ray burst GRB~980425, is discussed in particular detail.  Its radio properties are compared and contrasted with those of other known RSNe. ",
        "introduction": "A series of papers published over the past 20 years on radio supernovae (RSNe) has established the radio detection and/or radio evolution for approximately two dozen objects: 2 Type Ib supernovae (SNe), 5 Type Ic SNe, and the rest Type II SNe.  A much larger list of more than 100 additional SNe have low radio upper limits (Table 1).  \\begin{table} \\begin{center} \\scriptsize \\begin{tabular}{lcclcclcclcc} SN & Type & Radio & SN & Type & Radio & SN & Type & Radio & SN & Type & Radio \\\\ \\hline 1895B  & I     &    & 1901B  & I     &    & 1909A  & II    &    & 1914A  & ?     &    \\\\ 1917A  & I?    &    & 1921B  & II    &    & 1921C  & I     &    & 1923A  & IIP   & DT \\\\ 1937C  & Ia    &    & 1937F  & II    &    & 1939C  & II    &    & 1940A  & IIL   &    \\\\ 1945B  & ?     &    & 1948B  & II?   &    & 1950B  & II?   & DT & 1954A  & I     &    \\\\ 1954J  & II    &    & 1957D  & II?   & DT & 1959D  & II    &    & 1959E  & I     &    \\\\ 1963J  & I     &    & 1966B  & II    &    & 1968D  & II    & DT & 1968L  & IIP   &    \\\\ 1969L  & IIP   &    & 1970A  & II?   &    & 1970G  & IIL   & LC & 1970L  & I?    &    \\\\ 1970O  & ?     &    & 1971G  & I?    &    & 1971I  & Ia    &    & 1971L  & Ia    &    \\\\ 1972E  & Ia    &    & 1973R  & IIP   &    & 1974E  & ?     &    & 1974G  & Ia    &    \\\\ 1975N  & Ia    &    & 1977B  & ?     &    & 1978B  & II    &    & 1978G  & II    &    \\\\ 1978K  & II    & LC & 1979B  & Ia    &    & 1979C  & IIL   & LC & 1980D  & IIP   &    \\\\ 1980I  & Ia    &    & 1980K  & IIL   & LC & 1980L  & ?     &    & 1980N  & Ia    &    \\\\ 1980O  & II    &    & 1981A  & II    &    & 1981B  & Ia    &    & 1981K  & II?   & LC \\\\ 1982E  & Ia?   &    & 1982R  & Ib?   &    & 1983G  & I     &    & 1983K  & IIP/n &    \\\\ 1983N  & Ib    & LC & 1984A  & I     &    & 1984E  & IIL/n &    & 1984L  & Ib    & LC \\\\ 1984R  & ?     &    & 1985A  & Ia    &    & 1985B  & Ia    &    & 1985F  & Ib    &    \\\\ 1985G  & IIP   &    & 1985H  & II    &    & 1985L  & IIL   & DT & 1986A  & Ia    &    \\\\ 1986E  & IIL   & LC & 1986G  & Ia    &    & 1986I  & IIP   &    & 1986J  & IIn   & LC \\\\ 1986O  & Ia    &    & 1987A  & IIpec & LC & 1987B  & IIn   &    & 1987D  & Ia    &    \\\\ 1987F  & IIpec &    & 1987K  & IIb   &    & 1987M  & Ic    &    & 1987N  & Ia    &    \\\\ 1988I  & IIn   &    & 1988Z  & IIn   & LC & 1989B  & Ia    &    & 1989C  & IIn   &    \\\\ 1989L  & IIL   &    & 1989M  & Ia    &    & 1989R  & IIn   &    & 1990B  & Ic    & LC \\\\ 1990K  & IIL   &    & 1990M  & Ia    &    & 1991T  & Ia    &    & 1991ae & IIn   &    \\\\ 1991ar & Ic    &    & 1991av & IIn   &    & 1991bg & Ia    &    & 1992A  & Ia    &    \\\\ 1992H  & IIP   &    & 1992ad & IIP?  & DT & 1992bd & II    &    & 1993G  & II    &    \\\\ 1993J  & IIb   & LC & 1993N  & IIn   &    & 1993X  & II    &    & 1994D  & Ia    &    \\\\ 1994I  & Ic    & LC & 1994P  & II    &    & 1994W  & IIn   &    & 1994Y  & IIn   &    \\\\ 1994ai & Ic    &    & 1994ak & IIn   &    & 1995G  & IIn   &    & 1995N  & IIn   & DT \\\\ 1995X  & IIP   &    & 1995ad & IIP   &    & 1995al & Ia    &    & 1996N  & Ic    & DT \\\\ 1996W  & IIpec &    & 1996X  & Ia    &    & 1996ae & IIn   &    & 1996an & II    &    \\\\ 1996aq & Ic    &    & 1996bu & IIn   &    & 1996cb & IIb?  & DT & 1997X  & Ic    & DT \\\\ 1998bw & Ic    & LC \\\\ \\end{tabular} \\end{center} \\caption{Observed Supernovae (DT = Detection; LC = Light Curve Available)} \\end{table}  In this extensive study of the radio emission from SNe, several effects have been noted: 1) Type Ia SNe are not radio emitters to the detection limit of the VLA\\footnote{The VLA is operated by the NRAO of the AUI under a cooperative agreement with the NSF.}; 2) Type Ib/c SNe are radio luminous with steep spectral indices (generally $\\alpha < -1$; $S \\propto \\nu^{+\\alpha}$) and a fast turn-on/turn-off, usually peaking at 6 cm near or before optical maximum; and 3) Type II SNe show a range of radio luminosities with flat spectral indices (generally $\\alpha > -1$) and a relatively slow turn-on/turn-off, usually peaking at 6 cm significantly after optical maximum. Type Ib/c may be fairly homogeneous in their radio properties, while Type II, as in the optical, are quite diverse.  There are a large number of physical properties of SNe which we can determine from radio observations.  VLBI imaging shows the symmetry of the explosion and the local CSM, estimates the speed and deceleration of the SN shock propagating outward from the explosion and, with assumptions of symmetry and optical line/radio sphere velocities, allows independent distance estimates to be made (see, \\eg \\cite{Marcaide97}, \\cite{Bartel85}).  Measurements of the multi-frequency radio light curves and their evolution with time show the density and structure of the CSM, evidence for possible binary companions, clumpiness or filamentation in the presupernova wind, mass-loss rates and changes therein for the presupernova stellar system and, through stellar evolution models, estimates of the ZAMS presupernova stellar mass and the stages through which the star passed on its way to explosion.  It has also been proposed by Weiler \\etal\\/(1998) that the time from explosion to 6 cm radio maximum may be an indicator of the radio luminosity and thus an independent distance indicator.  A summary of the radio information on SNe can be found at {\\it http://rsd-www.nrl.navy.\\\\mil/7214/weiler/sne-home.html}. ",
        "conclusions": "Arising from one of the most energetic phenomena in the Universe, the radio emission from supernovae appears to be relatively well understood in terms of shock interaction with a structured circumstellar medium as described by the Chevalier (1982a,b) model and its modifications (\\cite{Weiler86}, \\cite{Weiler90}, and \\cite{Weiler91}). With this modeling, the radio emission can be used to estimate the circumstellar density, the pre-SN mass loss rate and changes therein, to show the existence of filamentation in the pre-SN stellar wind, to indicate the possible presence of binary companions, and to measure the symmetry of the explosion and the CSM with which it is interacting.  More speculatively, radio observations may also lead to a more physical classification system for SNe and provide a new technique for estimating distances to SNe and their parent galaxies.  However, the recent suggestion of an association of the $\\gamma$-ray burst GRB~980425 with the Type Ic supernova SN 1998bw provides evidence for yet a new phenomenon which may arise in at least some types of SN explosions.  Since SN 1998bw is a strong radio emitter, it is important to consider how it compares with the known properties of RSNe discussed above."
    },
    "0002/astro-ph0002051_arXiv.txt": {
        "abstract": "\\noindent We have used the {\\it ROSAT} PSPC to study the properties of a sample of 24 X-ray bright galaxy groups, representing the largest sample examined in detail to date. Hot plasma models are fitted to the spectral data to derive temperatures, and modified King models are used to characterise the surface brightness profiles.  In agreement with previous work, we find evidence for the presence of two components in the surface brightness profiles.  The extended component is generally found to be much flatter than that observed in galaxy clusters, and there is evidence that the profiles follow a trend with system mass. We derive relationships between X-ray luminosity, temperature and optical velocity dispersion. The relation between X-ray luminosity and temperature is found to be $L_X \\propto T^{4.9}$, which is significantly steeper than the same relation in galaxy clusters. These results are in good agreement with preheating models, in which galaxy winds raise the internal energy of the gas, inhibiting its collapse into the shallow potential wells of poor systems. ",
        "introduction": "\\label{sec:intro}  The majority of galaxies in the universe are found in galaxy groups \\cite{tully87}. These collections of between 3 and about 30 galaxies trace large scale structure (Ramella, Geller \\& Huchra 1990)\\nocite{ramella90} and probably contain a large fraction of the total baryonic mass in the universe (Fukugita, Hogan \\& Peebles 1998)\\nocite{fukugita98}.  However despite their abundance and importance, galaxy groups have received relatively little attention until recently. The main problem has been the identification of the groups themselves. Even when redshift information is available, it is difficult to identify whether a group is truly bound, due to the problems of small number statistics and chance superpositions. In contrast, galaxy clusters which are easier to identify due to the larger number of members, have been extensively studied.  The detection of extended X-ray emission from hot gas in the group potential well provides the best evidence that a group is truly gravitationally bound. The study of this hot intragroup gas can provide important insights into the evolution and dynamics of the group and its member galaxies. Samples of X-ray bright groups were originally studied using the {\\it Einstein} satellite (e.g. \\pcite{price91}), but the introduction of the {\\it ROSAT} satellite with its improved sensitivity and resolution, allowed a more thorough analysis of these systems. Since the {\\it ROSAT} PSPC was first used to study X-ray bright groups \\cite{mulchaey93,ponman93a} a number of collections of groups have been studied (e.g. \\pcite{doe95,ponman96a,burns96,mulchaey96,mahdavi97,mulchaey98b}). However, none of these studies provides a uniform, detailed analysis of a reasonable sized sample of groups, based on high quality data. The largest samples have all been based on {\\it ROSAT} All Sky Survey (RASS) data, in which case properties other than the luminosities are difficult to determine, due to poor statistics resulting from the short exposures.  The study of \\scite{ponman96a} used a mixture of RASS and pointed data and identified 22 X-ray bright groups. These were all compact groups from the catalogue of \\scite{hickson82}. Such compact groups have the advantage that they can be easily identified on the sky due to the high projected over-densities of galaxies within them, but may be unrepresentative of groups as a whole.  The X-ray properties of these Hickson compact groups (HCGs) showed systematic departures from those of clusters, leading to the suggestion that they might be displaying the marks of energy injection into the intergalactic medium due to galaxy winds.  \\scite{mulchaey98b} (henceforth MZ98) used pointed PSPC data to study groups of both loose and compact morphology.  With a sample of only nine groups they were unable to derive reliable statistical results, however they found that properties such as the $L:T$ relation and surface brightness slope were indistinguishable from those of clusters, in contradiction to the results of \\scite{ponman96a}. If this is true, it suggests a fundamental difference between the properties of loose and compact groups. The main aim of the present work is to establish the X-ray properties of loose groups by means of a careful and uniform study of a larger sample of systems, and to establish whether they do, in fact, differ from compact groups in their X-ray properties. To allow direct comparison with the results of MZ98, their sample has been included within ours.  In \\S~\\ref{sec:sample_selection} we describe the sample selection and initial identification of the X-ray bright groups. The spectral and spatial analyses of the X-ray emission are described in \\S~\\ref{sec:spectral_analysis} and \\S~\\ref{sec:spatial_analysis}.  Results of the analysis, including correlations between the derived parameters, are presented in \\S~\\ref{sec:Distributionsandcorrelations}.  These results are compared with those of MZ98 in \\S~\\ref{Comparison}, and discussed in \\S~\\ref{Discussion}. Finally, our conclusions are summarized in \\S~\\ref{Conclusions}. Throughout this paper we use H$_0$~=~50~km~s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": ""
    },
    "0002/astro-ph0002098_arXiv.txt": {
        "abstract": "Recent observational studies of intermediate- and high-mass star-forming regions at submillimeter and infrared wavelengths are reviewed, and chemical diagnostics of the different physical components associated with young stellar objects are summarized. Procedures for determining the temperature, density and abundance profiles in the envelopes are outlined. A detailed study of a set of infrared-bright massive young stars reveals systematic increases in the gas/solid ratios, the abundances of evaporated molecules, and the fraction of heated ices with increasing temperature. Since these diverse phenomena involve a range of temperatures from $<100\\,$K to 1000~K, the enhanced temperatures must be communicated to both the inner and outer parts of the envelopes. This `global heating' plausibly results from the gradual dispersion of the envelopes with time. Similarities and differences with low-mass YSOs are discussed. The availability of accurate physical models will allow chemical models of ice evaporation followed by `hot core' chemistry to be tested in detail. ",
        "introduction": "Massive star-forming regions have traditionally been prime targets for astrochemistry owing to their bright molecular lines (e.g., Johansson et al.\\ 1984, Cummins et al.\\ 1986, Irvine et al.\\ 1987, Ohishi 1997). Massive young stellar objects (YSOs) have luminosities of $\\sim 10^4 - 10^6$ L$_\\odot$ and involve young O- and B-type stars.  Because their formation proceeds more rapidly than that of low-mass stars and involves ionizing radiation, substantial chemical differences may be expected. The formation of high mass stars is much less well understood than that of low-mass stars. For example, observational phenomena such as ultracompact H II regions, hot cores, masers and outflows have not yet been linked into a single evolutionary picture.  Chemistry may well be an important diagnostic tool in establishing such a sequence.  Most of the early work on massive star-forming regions has centered on two sources, Orion--KL and SgrB2. Numerous line surveys at millimeter (e.g., Blake et al.\\ 1987, Turner 1991) and submillimeter (Jewell et al.\\ 1989, Sutton et al.\\ 1991, 1995, Schilke et al.\\ 1997) wavelengths have led to an extensive inventory of molecules through identification of thousands of lines. In addition, the surveys have shown strong chemical variations between different sources.  In recent years, new observational tools have allowed a more detailed and systematic study of the envelopes of massive YSOs.  Submillimeter observations routinely sample smaller beams (typically 15$''$ vs.\\ 30$''$--1$'$ ) and higher critical densities ($\\geq 10^6$ vs.\\ $10^4$ cm$^{-3}$) than the earlier work.  Moreover, interferometers at 3 and 1 millimeter provide maps with resolutions of $0.5''$--5$''$. Finally, ground- and space-based infrared observations allow both the gas and the ices to be sampled (e.g., Evans et al.\\ 1991, van Dishoeck et al.\\ 1999). These observational developments have led to a revival of the study of massive star formation within the last few years.  Recent overviews of the physical aspects of high-mass star formation are found in Churchwell (1999) and Garay \\& Lizano (1999).   In this brief review, we will first summarize available observational diagnostics to study the different phases and physical components associated with massive star formation.  Subsequently, an overview of recent results on intermediate mass YSOs is given, which are often better characterized than their high-mass counterparts because of their closer distance. Subsequently, we will discuss a specific sample of embedded massive YSOs which have been studied through a combination of infrared and submillimeter data. After illustrating the modeling techniques, we address the question how the observed chemical variations are related to evolutionary effects, different conditions in the envelope (e.g., $T$, mass) or different luminosities of the YSOs.  More extensive overviews of the chemical evolution of star-forming regions are given by van Dishoeck \\& Blake (1998), Hartquist et al.\\ (1998), van Dishoeck \\& Hogerheijde (1999) and Langer et al.\\ (2000).  Schilke et al.\\ (this volume) present high spatial resolution interferometer studies, whereas Macdonald \\& Thompson (this volume) focus on submillimeter data of hot core/ultracompact H~II regions. Ices are discussed by Ehrenfreund \\& Schutte (this volume).  \\begin{figure}[t] \\plotfiddle{dishoeck_fig1.ps}{11.0cm}{0}{50}{50}{-150}{-30} \\caption{Overview of the SEST 230 GHz line survey of SgrB2 at the N, M and NW positions, illustrating the chemical differentiation in high-mass star-forming regions (from: Nummelin et al.\\ 1998).} \\label{fig-1} \\vspace{-0.3cm} \\end{figure} ",
        "conclusions": ""
    },
    "0002/astro-ph0002117_arXiv.txt": {
        "abstract": "We report the status of a search for pulsars in the Galactic Centre, using a completely revised and improved high-sensitivity double-horn system at 4.85-GHz. We also present calculations about the success rate of periodicity searches for such a survey, showing that in contrast to conclusions in recent literature pulsars can be indeed detected at the chosen search frequency. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002321_arXiv.txt": {
        "abstract": "Two absorption line indices, Ca II and H$\\delta$/$\\lambda$4045, measured from high-resolution spectra are used with evolutionary synthesis models to verify the presence of a young stellar population in NGC~5018.  The derived age of this population is $\\sim$ 2.8 Gyr with a metallicity roughly solar and it completely dominates the integrated light of the galaxy near 4000 \\AA. ",
        "introduction": "Morphological peculiarities in the optical images of galaxies are now almost invariably taken to be signs of a past tidal interaction or merger.  Though theoretical models have had success reproducing tidal tails, shells, and other structures \\citep[e.g.,][]{tt72,q84,hq88,mbr93}, in the absence of explicit details of the interaction such as the Hubble types of the progenitors, relative progenitor sizes, and impact parameter, a unique solution is difficult to come by.  Dynamical friction during a collision almost certainly leads to a merger of the stellar systems \\citep{s83} and simulations have shown that any accompanying gas will rapidly dissipate to the center during minor mergers involving either ellipticals \\citep{wh93} or disk galaxies \\citep{mh94a}.  If the conditions are right, it is reasonable to expect that this gas inflow will result in star formation.  The duration, intensity, and even the starting time of the star formation, however, depend on the morphology of the progenitor and the details of the interaction \\citep{mh94b}.  As much post-merger information as possible is needed to help constrain the merger possibilities, including analysis of the resulting young stellar population (YSP).  Unfortunately, unless the merger remnant is presently forming stars at a reasonably vigorous rate (hence producing readily apparent emission lines) or has experienced a very recent and relatively strong episode of star formation placing it in the starburst regime, a YSP can be difficult to detect.  Broadband colors and other YSP indicators return to pre-star formation levels very quickly \\citep{bas90,cs94} and many indicators that imply the presence of a YSP can also be explained by intrinsic metallicity differences in the final population of the merger remnants \\citep[hereafter BBB]{bbb93}.  Ambiguity about whether a YSP even exists or not complicates the details of the merger considerably.  In this paper, we use spectral indices in conjunction with evolutionary synthesis models to detect and determine the age and metallicity of a YSP in the particular case of a possible merger remnant, NGC~5018.  It represents an update to the age-dating technique introduced in \\citet[hereafter LR]{lr96}.  NGC~5018 is appropriate to the present discussion because there is ongoing uncertainty concerning the presence of a YSP.  Certain observations imply the existence of a YSP while others seem to be inconsistent with the presence of a YSP.  NGC~5018 and the roots of the controversy are described in \\S2.  In \\S3, a review of the age-dating technique and its refinements are given.  The results of the technique applied to NGC~5018 are presented in \\S4 and \\S5 contains the conclusions. ",
        "conclusions": "The two observations of BBB which led them to conclude that a YSP is not present in NGC~5018 are a low $\\mathrm{Mg_{2}}$ index coupled with the lack of an upturn in the UV SED both comparable to that found in M32. As seen in Figures~\\ref{fig:yspm17}--\\ref{fig:yspp00}, however, NGC~5018 lies in a different region of the H$\\delta$/$\\lambda$4045--Ca II diagram than M32 hence it cannot be just a high luminosity version of M32.  With the patchiness of the dust in the central regions of NGC~5018, there may be as much as two magnitudes of extinction in the far-UV part of the spectrum which could explain the lack of an upturn \\citep{cd94}.  The results derived here, however, indicate that large amounts of dust obscuration need not be invoked, although a modest amount of extinction would bring Figure \\ref{fig:lick} into better harmony with the spectral index results.  The YSP age of~2.8 Gyr derived here is reddening insensitive and thus robust regardless of the structure of the dust distribution.  Also, as Figure~\\ref{fig:yspp00} shows, the YSP is completely dominating the visible part of the spectrum of NGC~5018.  A YSP of this age does not have an upturn in the UV part of the spectrum. Hence the small far-UV upturn in NGC~5018 may simply reflect the characteristics of the 2.8-Gyr-old stellar population dominating the light.  The models predict that a YSP of this age is about 6 times brighter (at 4000\\AA ) than an ancient population of the same metallicity and IMF. Hence if the relative contributions from the old and young populations were only half and half, then the galaxy must be 6 parts old to 1 part young in the central region.  This would qualify as a gas-rich merger event with a rapid gas inflow to the center \\citep{wh93} of NGC 5018.  If the light contribution is weighted even more heavily toward the young population as we suggest, this implies a more violent event, perhaps a merger of two spiral galaxies or some other event involving large amounts of gas consumed in star formation, with an extremely high inflow to the nucleus. Note that, if nebular emission is filling in the Balmer lines, the derived age decreases, the YSP is brighter, and so the size of the merging event decreases.  A solar-[Fe/H] population of age 2.8 Gyr has an Mg$_2 \\approx 0.20 $ mag, as is observed in NGC~5018. Aging this population to 15 Gyr raises the Mg$_2$ to 0.27 mag, still about 2 standard deviations from the mean relation for elliptical galaxies, though a dust-hidden YSP could be diluting the Mg$_2$ index \\citep{cd94}. We note, however, that compared to M32, NGC~5018's C$_2$4668 index is weaker but its Mg$_2$ index is stronger, implying that NGC~5018 may also participate in the general trend for large ellipticals to have enhanced Mg abundance \\citep[e.g.,][]{w98}. This will further increase NGC~5018's Mg$_2$ line strength as it ages so that it may one day fall among other ellipticals in the Mg$_2$-$\\sigma$ relation.  As mentioned previously, significant emission has been observed in the center of NGC~5018 \\citep{ghjn94}.  Since both of the spectral indices we used depend on the intensity in the center of a Balmer line, contamination by emission could affect the results significantly even though the contamination decreases as one moves toward higher order lines in the Balmer sequence.  Specifically, emission will weaken H$\\delta$ relative to Fe I $\\lambda$4045 and to a lesser degree weaken Ca II H+H$\\epsilon$ relative to Ca II K.  Each index, if contaminated, will therefore have a higher value than if there were no emission.  From Figures~\\ref{fig:yspm17}--\\ref{fig:yspp00}, it can be seen that if emission contamination were removed, the data points for NGC~5018 would shift toward younger ages.  The determined YSP age of 2.8 Gyr represents therefore an \\emph{upper limit on the YSP age}.  The limited spectral coverage does not permit a more detailed analysis of the emission contamination, but the fact that H$\\beta$ from Figure \\ref{fig:lick} gives the same age indicates that the emission must be relatively modest.  The power of the method used in this paper is its ability to discriminate between different, plausible stellar populations.  The verification of the presence of a YSP in NGC~5018 and its age of~2.8 Gyr are quite unambiguous, irrespective of the nature of the old, underlying population.  Also, this type of determination is not unique to NGC~5018.  A similar procedure can be applied to any galaxy with morphological peculiarities suspected of harboring more than one coeval stellar population.  With this tool, the effects of a dynamical interaction on the integrated light of a galaxy can be analyzed more fully."
    },
    "0002/hep-ph0002223_arXiv.txt": {
        "abstract": "Constraints on the lifetime of the heavy sterile neutrino, that was proposed as a possible interpretation of the KARMEN anomaly, are derived from primordial nucleosynthesis and SN~1987A. Together with the recent experimental bounds on the $\\nus$ lifetime, SN~1987A completely excludes this interpretation. Nucleosynthesis arguments permit a narrow window for the lifetime in the interval 0.1--0.2~sec. If $\\nus$ possesses an anomalous interaction with nucleons, the SN bounds may not apply, while the nucleosynthesis ones would remain valid. ",
        "introduction": "In 1995 the KARMEN collaboration discovered an anomaly in the time distribution of the charged and neutral current events induced by neutrinos from $\\pi^+$ and $\\mu^+$ decays at rest~\\cite{karmen95}. This anomaly may be explained by the production of a new neutral particle in pion decay \\be \\pi^+ \\rar \\mu^+ + x^0~, \\label{piondecay} \\ee with the mass 33.9 MeV, barely permitted by the phase space, so that this particle moves with non-relativistic velocity. Its subsequent neutrino-producing decays could be the source of the delayed neutrinos observed in the experiment. The anomaly was recently confirmed by the same group~\\cite{karmen99} with better statistics and substantially reduced cosmic-ray background.  A search for $x^0$ particles produced in the rare pion decay (\\ref{piondecay}) was performed in 1995 by an experiment at PSI~\\cite{psi95}. It gave an upper limit for the branching ratio of ${\\rm BR}( \\pi^+ \\rar \\mu^+ + x^0) < 2.6 \\times 10^{-8}$ at 95\\% CL.  Several candidates for $x^0$ have been proposed in the literature.  In ref.~\\cite{barger95} the authors considered a sterile neutrino, $x^0 \\equiv \\nu_s$.  Their conclusion was that the sterile-neutrino hypothesis is compatible with all laboratory constraints, but they noted possible problems with astrophysics and cosmology.  Further laboratory constraints on this model were investigated in ref.~\\cite{govaerts96}, concluding that mixings of $\\nu_s$ with $\\nue$ and $\\num$ must be very small, while a mixing with $\\nut$ was permitted. In this case $\\nu_s$ would predominantly decay through neutral-current interactions into $\\nut + \\ell + \\bar\\ell$, where $\\ell$ is any light lepton, $\\ell=\\nue$, $\\num$, $\\nut$, or $e^-$. The lifetime of $\\nu_s$ with respect to this decay was estimated to be in the range \\be 10^{-3}~{\\rm sec}<\\tau_{\\nu_s} < 150~{\\rm sec}~. \\label{taunus} \\ee The lower bound on $\\tau_{\\nu_s}$, which comes from the experimental bound on the $\\nu_\\tau$ mass (large mixing with $\\nu_s$ makes it too heavy), was discussed in ref.~\\cite{ahluwalia97}.  In ref.~\\cite{choudhury96} it was suggested that $x^0\\equiv\\tilde\\chi$ could be the lightest supersymmetric particle, photino or zino, that decayed through the channel \\be \\tilde \\chi \\rar \\gamma + \\num~. \\label{chi2} \\ee The new data~\\cite{karmen99}, however, do not agree with the predictions of the model so that recently a new version of supersymmetric model was considered~\\cite{choudhury99}, according to which the light neutralino decayed through a three body channel \\be \\tilde \\chi \\rar e^++e^-+\\nu_{\\mu,\\tau}~. \\label{chi3} \\ee  In the paper~\\cite{gninenko98} a new decay mode of muons was proposed as a source of the anomaly: \\be \\mu^+ \\rar e^+ + S~, \\label{scalar} \\ee where $S$ is a scalar boson with mass 103.9~MeV. A search for these decays was reported in ref.~\\cite{bilger99} where the upper limit ${\\rm BR}(\\mu^+\\rar e^+ + S) < 5.7\\times 10^{-4}$ was obtained. This limit, though it does not exclude the model, makes it more complicated.  Recent data by the NOMAD collaboration~\\cite{nomad} permit to strengthen the bound on the mixing of $\\nu_s$ with $\\nu_\\tau$. Their expected lower limit on $\\tau_{\\nu_s}$ is around 0.1~sec.  On the other hand, cosmology and astrophysics permit to obtain an upper bound on $\\tau_{\\nu_s}$ that may be complementary to direct experiments. If this happens to be the case, then the explanation of the KARMEN anomaly by a 33.9 MeV sterile neutrino would be ruled out. Recently two papers have appeared~\\cite{goldman99,kach99} where the observation of SN~1987A were used to put bounds on the properties of the proposed sterile neutrino or light neutralino. In the present paper we derive constraints on the lifetime of a 33.9~MeV sterile neutrino from both more detailed consideration of SN~1987A and of primordial nucleosynthesis. With the present-day accuracy of the data on the primordial light-element abundances, the bounds that are found from SN~1987A tend to be stronger than those found from primordial nucleosynthesis. However, the latter remain interesting, first, because they may be competitive in the nearest future with an improved accuracy of the data on primordial nucleosynthesis, and second, in a hypothetical case where $\\nus$ would have an anomalously strong (stronger than the usual weak) interaction with nucleons. In that case the SN~1987A bounds would not apply, while the nucleosynthesis bounds would survive. ",
        "conclusions": "Cosmological and astrophysical arguments seem to exclude the interpretation of the KARMEN anomaly by an unstable sterile neutrino mixed with $\\nut$.  The arguments based on SN~1987A are stronger than the nucleosynthesis bound. The supernova limit is given by (\\ref{eq:mixinglimit}), i.e.~$\\taus > 6\\times 10^4$ sec.  This result, together with the direct experimental constraints on $\\nut$, completely excludes a 33.9~MeV sterile neutrino.  Primordial nucleosynthesis permits to exclude roughly $\\taus >0.2 $ sec.  So with the existing direct experimental limits, some window, $\\taus = 0.1-0.2$ sec, remains open.  The simplifications that we made in deriving the BBN bound typically lead to a weaker result, so the real bound may be somewhat stronger. However, if exact calculations confirm the decrease of ${}^4$He around $\\taus = 0.1$ found in this paper, then BBN will never exclude lifetimes $\\sim 0.1$ sec.  The effect happens to be surprisingly sensitive to usually neglected phenomena, particularly the shape of the $\\nue$ spectrum and to the non-adiabatic variation of temperature.  We hope to do exact (but rather long) calculations later. Together with a possible improvement of the observational data on light element abundances and a better understanding of the subsequent changes of light elements in the course of cosmological evolution, the BBN bound may become competitive with the supernova one. If so, the very attractive hypothesis of a 33.9~MeV sterile $\\nus$ would be killed by two independent arguments. Moreover, one may imagine a case when the supernova arguments are not applicable, while the nucleosynthesis ones still operate. For example, if $\\nus$ possesses an anomalously strong interaction with nucleons (stronger than the usual weak one), then it would not noticeably change the energetics of supernovae that we described here, but would affect nucleosynthesis practically at the same level as discussed in section~\\ref{ns}.  This new interaction might be related to the anomalously high mass of $\\nus$. If this is the case then a small window for a 33.9~MeV $\\nus$ with lifetime $\\tau_{\\nu_s} = 0.1$--0.2~sec, still exists at the present time, and stronger experimental bounds, as well as more accurate calculations of the impact of $\\nus$ on primordial nucleosynthesis, are needed."
    },
    "0002/astro-ph0002185_arXiv.txt": {
        "abstract": "The K-band Hubble diagram of Brightest Cluster Galaxies (BCGs) is presented for a large, \\xray\\ selected cluster sample extending out to $z = 0.8$. The controversy over the degree of BCG evolution is shown to be due to sample selection, since the BCG luminosity depends upon the cluster environment. Selecting only the most \\xray\\ luminous clusters produces a BCG sample which shows, under the assumption of an Einstein-de Sitter cosmology, significantly less mass growth than that predicted by current \\semianalytic\\ galaxy formation models, and significant evidence of any growth only if the dominant stellar population of the BCGs formed relatively recently ($z \\leq 2.6$). ",
        "introduction": "\\label{intro}  The majority of stars in giant ellipticals found in the cores of rich galaxy clusters are old; photometric and spectroscopic studies of cluster galaxies out to $z \\approx 1$ suggest a formation redshift, \\zf, greater than 2, with little variation within a cluster, and that secondary bursts of star formation account for a small fraction of the stellar mass (\\eg\\ Arag\\'{o}n-Salamanca \\etal~\\markcite{aecc93}1993; Ellis \\etal~\\markcite{ellis97}1997; Stanford, Eisenhardt \\& Dickinson~\\markcite{sed98}1998; van Dokkum \\etal~\\markcite{vd98}1998; Poggianti \\etal~\\markcite{morphs99}1999).  However, to understand the process of galaxy formation it is necessary to know where the stars were formed as well as when. In the traditional view of early-type galaxy formation---a ``monolithic'' collapse at high redshift (\\eg\\ Eggen, Lynden-Bell \\& Sandage~\\markcite{els62}1962; Larson~\\markcite{larson69}1969)---all the stars were formed in situ, in direct contrast to the merger-driven growth of galaxies predicted by \\semianalytical\\ models for hierarchical cosmologies, such as CDM (\\eg\\ Kauffmann \\& White~\\markcite{kw93}1993; Baugh, Cole \\& Frenk~\\markcite{bcf96}1996). Since the ages of the stars are similar in both scenarios, it is the change in mass with look-back time that separates the two pictures observationally. The current data are inconclusive; for example, the hierarchical models are \\favoured\\ by the enhanced merger fraction seen in the $z=0.8$ cluster MS1054.4-0321 (van Dokkum \\etal~\\markcite{v99}1999), whilst \\deprop\\ show no evidence for mass evolution of bright ellipticals in clusters out to $z \\approx 1$. In general it is difficult to follow the evolutionary history of ellipticals since selection methods can seriously bias the samples, c.f. the discussions of progenitor bias in van Dokkum \\& Franx~\\markcite{vf96}(1996), the effect of preferential selection of the most massive objects at each epoch in Kauffmann \\& Charlot~\\markcite{kc98}(1998), and the use of \\colour\\ selection in Jimenez \\etal~\\markcite{jimenez99}(1999).  One approach to minimising such problems is to study the evolution of a particular class of ellipticals---brightest cluster galaxies (BCGs)---because of their unique location, close to the \\centre\\ of the cluster's potential well. BCGs do not appear to be drawn from the same luminosity function as other cluster galaxies (\\eg\\ Dressler~\\markcite{dressler78}1978), which suggests that they have a distinct formation history. Knowledge of BCG evolution can therefore provide different constraints on galaxy formation models to studies of the general cluster population.  The K-band Hubble diagram for BCGs has recently been extended to $z \\approx 1$ by both \\cmfirst\\ and \\abkfirst: these observations provide the opportunity to measure the luminosity evolution of BCGs since evolutionary and pass-band corrections are insensitive to the recent star-formation history of a galaxy at near-IR wavelengths (\\eg\\ Bershady~\\markcite{bershady95}1995; Madau, Pozzetti, \\& Dickinson~\\markcite{madau98}1998). The conclusions drawn are contradictory, despite the use of the same cosmology and a common assumption that the stellar populations of BCGs are old and passively evolving; \\cmmain\\ assert that the stellar populations of BCGs in the most massive clusters have not grown significantly since $z \\approx 1$, whilst \\abkmain\\ argue that their results are in good agreement with the mass increase of BCGs---by a factor of four in an Einstein-de Sitter cosmology---predicted by \\semianalytical\\ models over the same redshift range. The two samples have almost no overlap---\\cmmain\\ having used an \\xray\\ selected cluster catalogue whilst \\abkmain\\ used a heterogeneous compilation that was mainly optically selected---and it is the aim of the present work to show that the results can be reconciled by considering the properties of the clusters in the two samples. Section~\\ref{data} describes the BCG sample used--- an extension of that of \\cmmain---and the reduction methods employed, whilst section~\\ref{results} presents the results of the analysis and a comparison to those of \\abkmain. Throughout this letter an Einstein-de Sitter cosmology with \\ho{50} is assumed, and \\xray\\ luminosities (\\lx) are quoted for the 0.3--3.5~keV pass band. ",
        "conclusions": "\\label{conclusion}  The K-band luminosities of BCGs are correlated with their environment: clusters with a high \\xray\\ luminosity contain BCGs which are brighter, and have a smaller scatter, than those BCGs in clusters with a low \\xray\\ luminosity. The BCG evolution seen by \\abkmain\\ has been shown to be an artifact of a selection bias in their cluster sample; at high redshifts, their clusters are systematically less \\xray\\ luminous than their low-redshift sample, and so their BCGs are systematically fainter.  Under the assumption of an Einstein-de Sitter universe, non-parametric tests show that the only significant evidence for BCG mass evolution over the range $0.05 \\leq z \\leq 0.83$ occurs when the dominant stellar population formed relatively recently ($\\zf \\leq 2.6$). Using the same parametric form as \\abkmain, the masses of BCGs in high-\\lx\\ clusters are found to have, at most, doubled since $z=1$, compared to the factor of $\\sim 4$ increase predicted, for BCGs in massive clusters, by the \\semianalytic\\ models discussed by \\abkmain."
    },
    "0002/astro-ph0002466_arXiv.txt": {
        "abstract": "New photometric observations of seven intermediate mass pre-main sequence $\\delta$ Scuti candidates are presented. The periods and pulsation modes are derived for two of these stars, namely HD~35929 and V351~Ori. The comparison between observations and nonlinear pulsational models allows us to provide some initial constraints on their mass and evolutionary state. As an illustration we discuss the use of periods to identify the mode of pulsation in these two stars and to have an independent estimate of their distances. ",
        "introduction": "Young pre--main-sequence (PMS) stars of mass greater than $\\sim$1.5 M$_\\odot$ cross the region of pulsational instability during their contraction toward the main-sequence. The location of the instability strip in the H-R diagram has been recently identified by means of nonlinear models for the first three radial modes (Marconi \\& Palla 1998). The time spent by intermediate-mass stars within the boundaries of the strip represents a small fraction of the Kelvin-Helmoltz time scale, varying between $\\sim$10\\% for a star of 1.5 M$_\\odot$ and $\\sim$5\\% for a 4 M$_\\odot$ star. Despite the brevity of this phase, however, a number of known Herbig Ae stars have the appropriate combination of luminosity and effective temperature to become pulsationally unstable. In our previous study, we suggested to look for $\\delta$~Scuti-type photometric variations with periods of minutes to several hours and amplitudes less than few tenths of magnitudes in a sample of Herbig Ae stars whose position in the H-R diagram coincides with the instability strip. \\\\ The identification of a few PMS objects pulsating as $\\delta$ Scuti stars (Breger 1972), the prototype being the star HR~5999 (Kurtz \\& Marang 1995), has provided some support to the connection between variability and stellar pulsation. Of particular interest is the Ae star HD~104237 that shows both short- and long-term velocity changes of spectral lines (Donati et al. 1997). These variations indicate that the star is undergoing radial pulsations with a period of approximately 40 minutes and an amplitude of about 1 km s$^{-1}$ (see also B\\\"ohm et al. 2000, in preparation). Interestingly, HD~104237 is the first intermediate mass PMS star with a measured magnetic field (Donati et al. 1997). A few new PMS candidate pulsators have been recently identified by Pigulski et al. (2000).\\\\ Stimulated by these initial results, we have started a photometric investigation of a sample of seven Herbig Ae stars with spectral types in the range A5 to F5, located within or near the boundaries of the instability strip. For some of them, large time scale variations have been observed during the long term monitoring program of variable stars conducted at ESO (LTPV project: Sterken et al. 1995 and references therein). However, no information is available on their variability on time scales shorter than 2 or 3 days. \\\\ The main goal of our study is to detect and characterize the pulsation properties of young stars. This way, we can improve our knowledge of their internal structure and obtain unique constraints on the theoretical predictions of the models. Ultimately, the analysis of the pulsation characteristics can yield an indirect estimate of the stellar mass. This represents a powerful method for stars that are not part of the restricted group of spectroscopic binary systems. In this Letter, we report the discovery of two additional Herbig Ae stars which show evidence for variability of the $\\delta$ Scuti-type. ",
        "conclusions": ""
    },
    "0002/astro-ph0002520_arXiv.txt": {
        "abstract": " ",
        "introduction": "The 1990's will be remembered as a decade of discovery for cosmic microwave background (CMB) anisotropies.  The launch of the COBE satellite ushered in the decade in 1990 and lead to the first detection of CMB anisotropies at $>10^{\\circ}$ scales \\cite{COBE92}.  Through the decade, a combination of higher resolution experiments made the case for a rise in the anisotropy level on degree scales and a subsequent fall at arcminute scales \\cite{ScoSilWhi95}. The final year saw experiments, notably Toco and Boomerang, with sufficient angular resolution and sky coverage to localize a sharp peak in the anisotropy spectrum at approximately $0^{\\circ}\\!.5$ \\cite{TegZal00}.  In this review, we discuss the theoretical implications of these results and provide a roadmap for critical tests and uses of CMB anisotropies in the coming decade. ",
        "conclusions": "We are already well on on our way to extracting the cosmological information contained in the primary temperature anisotropies, specifically the angular diameter distance to recombination, the baryon density, the matter-radiation ratio at recombination, and the ``acausal'' (inflationary) nature and spectrum of the initial perturbations. Even if our simplest inflationary cold dark matter model is not correct in detail, these quantities will be measured in the next few years by long-duration ballooning, interferometry and the MAP satellite, {\\it if the acoustic nature of the peak at $\\ell\\sim 200$ is confirmed by the detection of a second peak}. In the long term,  the high sensitivity and wide frequency coverage of the Planck satellite and other future experiments should allow CMB polarization and secondary anisotropies to open new windows on the early universe and large-scale structure.  \\vskip 0.2truecm \\begin{quotation} \\footnotesize\\noindent {\\it Acknowledgements:} I would like to thank my collaborators through the years, especially N. Sugiyama and M. White, the organizers of RESCEU 1999, and the Yukawa Institute for its hospitality. This work was partially supported by NSF-9513835, the Keck Foundation, and a Sloan Fellowship. \\end{quotation}"
    },
    "0002/hep-ph0002129_arXiv.txt": {
        "abstract": "In presence of a magnetic field, photons can mix with any particle having a two--photon vertex. In theories with large compact extra--dimensions, there exists a hierachy of massive Kaluza--Klein gravitons that couple to any photon entering a magnetic field. We study this mixing and show that, in comparison with the four dimensional situation where the photon couples only to the massless graviton, the oscillation effect may be enhanced due to the existence of a large number of Kaluza--Klein modes.  We give the conditions for such an enhancement and then investigate the cosmological and astrophysical consequences of this phenomenon; we also discuss some laboratory experiments.  Axions also couple to photons in the same way; we discuss the effect of the existence of bulk axions in universes with large extra--dimensions. The results can also be applied to neutrino physics with extra--dimensions. ",
        "introduction": "It is well known \\cite{gertsenshtein62,zeldovich83} that photons can be converted into gravitons by a magnetic field in a standard four dimensional spacetime. The propagation eigenstates are then mixtures of photon and graviton interaction eigenstates. In quantum words, this mixing is due to the fact that any particle which has a two--photon vertex can be created by a photon entering an external electromagnetic field \\cite{adler71,raffelt88} and can oscillate coherently with the photon. Classically, this can be understood by the fact that an electromagnetic plane wave cannot radiate gravitationally in vacuum since its stress--energy tensor contains no quadrupole \\cite{landau}. But, a time varying quadrupole appears (due to interference) when an electromagnetic plane wave propagates through a constant magnetic field \\cite{gertsenshtein62,zeldovich83,magueijo94}. The implications of this effect on the cosmic microwave background (photons) has been considered \\cite{magueijo94,chen95} and it has been shown that it will be undetectable for standard cosmological magnetic fields \\cite{cillis96}. A similar effect also generically happens for axions (and for any particle having a two--photon vertex).  The photon--axion (see \\cite{kim87,raffelt90,turner90} for reviews on axions) mixing has yet been studied in details by many authors (see e.g. \\cite{raffelt88,sikivie83,sikivie84,morris}) and is used in experiments, since the pioneer work by Sikivie \\cite{sikivie83}, to put constraints on the axion parameters \\cite{bibber89,hagmann98,kim98,sikivie97} (see e.g. \\cite{raffelt98} for an up to date review on such experiments).  Recently, a lot of interest has been raised by models where the universe has large extra--dimensions \\cite{arkani98,arkani99}.  In such models, the Planck scale, $M_4$, is no longer a fundamental scale but is related to the fundamental mass scale of the $D$ dimensional theory, $M_D$, through \\cite{arkani99} \\begin{equation}\\label{planck} \\bar M_4^2\\equiv R^n M_D^{n+2}, \\end{equation} where $R$ is a length scale (usually taken to be the radius of the $n=D-4$ compact extra--dimensions) and $\\bar M_4\\equiv M_4/\\sqrt{8\\pi}$. $M_D$ can be significantly smaller than $\\bar M_4$ at the price of having large extra--dimensions. These ideas can be naturally embedded in fundamental string theories with a low string scale \\cite{witten96,lykken96,dienes98,antoniadis98} (see also \\cite{antoniadis90} for earlier discussions on TeV scale extra--dimensions).  In these models, gravity can propagate in the $D$ dimensional spacetime (bulk space time) whereas the standard model fields are localised on a 3--brane.  An effect of the compact extra--dimensions arises from interactions between the Kaluza--Klein (KK) excitations of the gravitons (or other bulk fields) which are seen in four dimensions as a tower of massive particles \\cite{giudice98,han98}.  Constraints on the size of these extra--dimensions can be obtained both from the laboratory physics \\cite{arkani99} and from astrophysics and cosmology (see e.g. \\cite{bdl}). For instance, the emission of KK gravitons induces an energy loss in many astrophysical objects \\cite{barger99} such as the Sun, red giants and supernovae SN1987A \\cite{cullen99} implying the lower bound $M_D>30-130\\,\\hbox{Tev}$ (2.1 - 9.2) TeV for the case of $n=2$ (3) extra--dimensions \\cite{barger99,cullen99}. Some authors \\cite{arkani98,antoniadis98,chang99,dudas2} also have recently investigated the possible presence of axions in the bulk which would be coupled to the brane degrees of freedom. Such a bulk axion also gives rise to a tower of KK-states as seen from a four dimensional point of view.  The goal of this article is to investigate the effects of the photon--KK graviton and photon--KK axion oscillations and to estimate their effects in cosmology and astrophysics, as well as terrestrial experiments. Since there is a large number of KK states with which the photon can mix, one can expect a departure from the usual four dimensional result. We first describe (\\S~\\ref{par2}) the photon--KK graviton system starting from a $D$ dimensional action, linearising it and compactifying it to four dimensions. We also show (\\S~\\ref{par5}) that the photon--axion mixing is described by the same formalism and lead to the same effects as for gravitons. Then, we turn to investigate the mixing in itself.  For that purpose, we sum up the known results of the mixing of a photon with a low mass particle in four dimensions (\\S~\\ref{par3_1}) and then discuss the most general case of a $D$ dimensional spacetime (\\S~\\ref{par3_2}).  The general expression for the oscillation probability is then evaluated in the particular cases of a five (\\S~\\ref{par3_3}) and of a six (\\S~\\ref{DD}) dimensional spacetime. We show that, as long as one is in a weak coupling regime, one can add the individual probabilities which leads to an enhancement of the oscillation probability if the KK modes are light enough. We also show that there exists a regime where the photon mixes strongly preferentially with a given KK mode.  We generalise our results to an inhomogeneous magnetic field (\\S~\\ref{par4}) when the probability of oscillation is small, and then turn to the cosmological and astrophysical situations where such effects may be observed.  We study the case of the cosmic microwave background (\\S~\\ref{par6_1}), of pulsars (\\S~\\ref{par6_2}) and of magnetars (\\S~\\ref{par5_3}). We show that even if the enhancement of the oscillation probability can be very important, it is still very difficult to observe this effect in known astrophysical systems. To finish (\\S~\\ref{par7}) we discuss laboratory experiments and particularly polarisation experiments. For that purpose, we describe the computation of the phase shift between the two polarisations of an electromagnetic wave and compare the result to the standard four dimensional case. ",
        "conclusions": "We have used the property that, in an external magnetic field, the KK gravitons and axions couple at tree level to photons to show that there exists a mixing between the photon and these particles. We have computed this mixing and compared our result to the photon-graviton and photon-axion mixings in four dimensions.  The main difference comes from the fact that the mixing matrix is now infinite and that a photon couples to a large number of massive particles.  We then have discussed the physical implications of these phenomena in a general $D$ dimensional universe.  This leads us to conclude that for most astrophysical objects the effect of photon-KK gravitons mixing will be unobservable.\\\\  The main points of this study are: \\begin{itemize} \\item We describe how to deal with the mixing between the photon and a large number of light particles. This extends the former results to the case of $D$ dimensional universes where a photon can mix with all the Kaluza--Klein gravitons and possibly  with bulk axions. In (\\ref{ppp}--\\ref{y}), we gave the exact expression for the oscillation probability and we then discussed its amplitudes first qualitatively and then in a five and in a six dimensional spacetime. \\item When $\\lambda_\\gamma<R$, i.e. if the caracteristic length scale associated with the photon wavelength and effective mass is smaller than the radius of the extra--dimensions, there is a departure from the four dimensional effect. Otherwise, the first KK mode is too heavy to be excited and everything, in general, reduces to the four dimensional situation. \\item Two limiting regimes have been found: \\begin{itemize} \\item A {\\it large radius regime}  where the two following behaviours can appear \\begin{itemize} \\item In the {\\it weak mixing} regime, we have shown that the oscillation probability can be obtained by summing over the individual oscillation probabilities and is then enhanced by a factor of order $\\beta^{-n/2}$ in comparison with the standard four dimensional case. In that case the solutions can be found either by solving the eigenvalues equation or by considering the equation of evolution as a Schr\\\"odinger equation and solving it iteratively in the interaction picture. The latter method generalises to the case of inhomogeneous magnetic fields but only applies if the oscillation probability is small compared to unity. \\item In the {\\it strong mixing} regime the photon mixes preferentially with a given KK modes, which is possible if the plasma effects dominate over the vacuum polarisation. In that case a complete transition is possible. We note that this effect is more likely to happen in a $D$ dimensional context than in a four dimensional spacetime and point out a specific effect of the extra--dimensions on the oscillation length. \\end{itemize} \\item A {\\it small radius regime} where $R\\ll\\lambda_\\gamma$ so that the spacing between to KK modes is very large compared to the photon characteristic length. In that case, the probability is generally dominated by the contribution of the first state corresponding to the lightest particle and we are back to the four dimensional case. A consequence of this is that we can have an observable signature of the extra--dimensions only if $\\lambda_\\gamma<R$. In most of the systems this cannot be achieved, mainly because $\\Delta_\\lambda$ is very small; the only favorable situation happens in strong magnetic fields such as in pulsars and magnetars and when we deal with a magnetic field which varies on a small enough typical scale. \\end{itemize} \\item We have shown that, in the case of graviton, the effect of this mixing although enhanced is too small to be observed on the cosmic microwave background and on astrophysical objects such as pulsars. However, we point out that the effects can be larger with bulk axions. \\item We discussed laboratory experiments designed for the search of axions in the light of this new framework and we computed the phase shift (\\ref{P1}--\\ref{P2}) between the two polarisations of a photon entering a magnetic field. As for the probability, we show that the phase shift is enhanced by the existence of the KK modes. In a periodic field, we show that the effect of the extra--dimensions can be important and that there exists a series of strong and weak mixing regimes that may be observed. More work to derive bounds on $f_{\\rm PQ}$ is however needed before drawing any conclusion. \\item On a more technical level, we have shown how to diagonalise a general matrix of mixing. This kind of matrices appears in different situations in extra-dimensions physics and these results can be used in many problems and in particular for neutrino oscillations. \\item To finish, let us stress some important comments. First, we have assumed the validity of the Euler--Heisenberg Lagrangian to describe the vacuum polarisation. One has also to be aware that in strong magnetic fields such as in pulsar magnetosphere, photon splitting \\cite{adler71,adler2} will be in competition with the photon graviton oscillation. We have not compared the strength of these two effects but we expect the latter to be dominant in high magnetic fields, such as in the magnetosphere of magnetars.  We have also concentrated on gravitational waves even if there is also production of scalar waves. They are thought to be negligible, at least in the four dimensional case \\cite{magueijo94}. Concerning the axions, the effects may be more important than for gravitons, depending on the value of the coupling $f_{\\rm PQ}$, but we did not reconsider the bounds on the axion parameters. Further work is needed in that direction. We also stress that in general one needs a precise determination of their mass spectrum to compute the coupling of each mode to the photon. Most of these results were obtained for $n=1$ or $n=2$ extra--dimensions for which the results respectively do not, or only weakly, depend on the cut--off of the theory. We also stress that depending on the exact physical situation the number of KK modes with which the photon can oscillate coherently can be drastically limited in comparison with the number of accessible KK modes. These decoherence effect implies that the UV cut--off of the theory $M_{\\rm max}$ is expected to be much higher than the physical cut--off. A consequence of this is that the results obtained in the cases $n=1$ and $n=2$ can be extended to $n>2$ without depending on $M_{\\rm max}$. \\end{itemize}"
    },
    "0002/astro-ph0002246_arXiv.txt": {
        "abstract": "The $I-Z$ colour has been recently shown to be a good temperature indicator for M dwarfs.  We present the first $IZ$ photometry of a small sample of L dwarfs ranging in spectral type from L0.5V to L6.0V.  We find that the $I-Z$ colour is not a good temperature indicator for objects between L1V and L5V, such objects having colours that overlap with mid M dwarfs. We attribute this to the reduction in the strength of the TiO and VO bands in the L dwarfs which are the dominant opacity source in the $I$ band for late M dwarfs. Beyond L5V, $I-Z$ appears to be a reasonable indicator.   This has important implications for the planning of optical surveys for cool objects in clusters and the field.  For example $I-Z$ will cease to be a good method of discriminating brown dwarfs in the Pleiades below around $0.04 M_\\odot$, and at around $0.075 M_\\odot$ in the Hyades and Praesepe. ",
        "introduction": "Kirkpatrick et al. (1999) have recently identified and classified the first large sample of L dwarfs (objects cooler than M-dwarfs) from the 2-micron all sky survey (2MASS).  The optical spectra of these objects are characterised by the disappearance of the TiO and VO bands which dominate late M dwarfs, and their replacement by metallic Hydrides and neutral alkali metals.  This would be expected to affect the optical colours one observes for such objects.  This is important as many surveys for low mass objects (specifically brown dwarfs) in clusters have relied on colours such as $V-I$ (e.g. Stauffer, Hamilton \\& Probst 1994), $R-I$ (e.g. Jameson \\& Skillen 1989, Hambly et al. 1999), and $I-Z$ (e.g  Pinfield et al. 1997, Cossburn et al. 1997, Zapatero Osorio et al. 1999).  The $I-Z$ colour has been particularly favoured of late, as it has been found to be an excellent method of picking late cluster M dwarfs from the field.  The lowest mass object so far detected in the Pleiades using this technique is the $\\sim$ L1V dwarf Roque 25 which was recently identified by Martin et al. (1998). It is therefore useful to see if the $I-Z$ colour would remain useful as a method of picking out later $L$ dwarfs in cluster fields.  For example at the distance of the Hyades an L3V dwarf would have $I\\sim 19$ based on the pseudo-photometry (derived from flux calibrated spectra) presented by Kirkpatrick et al. (1999). This is easily obtained with 2-m class telescopes, and therefore an optical survey at $I$ and $Z$, at which wavelengths much larger fields of view are generally available than the near infrared $JHK$ bands, would appear to be an excellent way of finding such objects.  In order to address this question of applicability to cluster surveys, as well as the more general question of the use of $I-Z$ as a temperature indicator we have carried out $I_{\\rm Harris}$ (hereafter $I_H$) and $Z_{RGO}$ photometry of a small number of L dwarfs ranging in spectral type from L0.5V to L6V.  This paper presents the results of that photometry, and discusses the results in the context set out above. ",
        "conclusions": "We have presented observations that show that the $I-Z$ colour is not a good temperature indicator for L dwarfs between L1V and L5V, these objects having $I-Z$ colours which overlap with mid-late M dwarfs.  We attribute this to the decreasing blanketing of the $I$ band flux in L dwarfs due to the absence of strong TiO and VO bands in their spectra.  This imposes limits on the use of $I-Z$ as an indicator of very cool objects in cluster brown dwarf searches at around $0.04 M_\\odot$ for the Pleiades and $0.075 M_\\odot$ for Praesepe and the Hyades.  For objects of lower mass than this near infrared ($JHK$) surveys should be considered."
    },
    "0002/astro-ph0002070_arXiv.txt": {
        "abstract": "\\noindent{\\it We present an overview of observational progress in the study of extragalactic globular cluster systems. Globular clusters turn out to be excellent tracers not only for the star-formation histories in galaxies, but also for kinematics at large galactocentric radii. Their properties can be used to efficiently constrain galaxy formation and evolution. After a brief introduction of the current methods and futures perspectives, we summarize the knowledge gained in various areas of galaxy research through the study of globular clusters. In particular, we address the star-formation histories of early-type galaxies; globular cluster population in late-type galaxies and their link to early-type galaxies; star and cluster formation during mergers and violent interactions; and the kinematics at large radii in early-type galaxies. The different points are reviewed within the context of galaxy formation and evolution.  Finally, we revisit the globular cluster luminosity function as a distance indicator. Despite its low popularity in the literature, we demonstrate that it ranks among one of the most precise distance indicators to early-type galaxies, provided that it is applied properly. } ",
        "introduction": "This brief review on extragalactic globular cluster systems is derived from a lecture given for the award of the Ludwig-Bierman-Preis of the Astronomische Gesellschaft in G\\\"ottingen during September 1999. The Oral version aimed at introducing, mostly from an observer's point of view, this field of research and at emphasizing its tight links to galaxy formation and evolution.  The scope of this written follow-up is {\\it not} to give a complete review on globular cluster systems but to present recent discoveries, including examples, and to set them into the context of galaxy formation and evolution. The choice of examples and the emphasis of certain ideas will necessarily be subjective, and we apologize at this point for any missing references.  Excellent recent reviews can be found in the form of two books: ``Globular Cluster Systems'' by Ashman \\& Zepf (1998), as well as ``Globular Cluster Systems'' by Harris (2000). These include a full description of the globular clusters in the Local group (not discussed here), as well as an extensive list of references, including to older reviews.  The plan of the article is the following. In section 2, we give an introduction and the motivation for studying globular cluster systems with the aim of understanding galaxies. Section 3 presents current and future methods of observations, and the rational behind them. This section reviews recent progress in optical and near-infrared photometry and multi-object, low-resolution spectroscopy. It can be skipped by readers interested in results rather than methods. In Section 4 we present in turn the status of our knowledge on globular cluster systems in ellipticals, spirals and mergers. What are the properties of the systems? How are the galaxy types linked? And do mergers produce real `globular' clusters? In section 5, we discuss sub-populations of globular clusters and their possible origin. The most popular scenarios to explain the presence of globular cluster sub-populations around galaxies are listed. The pros and cons, as well as the expectations of each scenario are discussed. We present, in section 6, some results from the study of globular cluster system kinematics. Finally, in section 7, we revisit the globular cluster luminosity function as a distance indicator. Under which conditions can it be used, and how should it be applied to minimize any systematic errors? It is compared to other distance indicators and shown to do very well. Some conclusions and an outlook are given in section 8. ",
        "conclusions": ""
    },
    "0002/astro-ph0002300_arXiv.txt": {
        "abstract": "We investigate the properties of $r$-mode and inertial mode of slowly rotating, non-isentropic, Newtonian stars, by taking account of the effects of the Coriolis force and the centrifugal force. The Coriolis force is the dominant restoring force for both $r$-mode and inertial mode, which are also called rotational mode in this paper. For the velocity field produced by the oscillation modes the $r$-mode has the dominant toroidal component over the spheroidal component, while the inertial mode has the comparable toroidal and spheroidal components. In non-isentropic stars the specific entropy of the fluid depends on the radial distance from the center, and the interior structure is in general divided into two kinds of layers of fluid stratification that is stable or unstable against convection. Because of the non-isentropic structure, low frequency oscillations of the star are affected by the buoyant force, which has no effects on oscillations of isentropic stars. In this paper we employ simple polytropic models with the polytropic index $n=1$ as the background neutron star models for the modal analysis. For the non-isentropic models we consider only two cases, that is, the models with the stable fluid stratification in the whole interior and the models that are fully convective. For simplicity we call these two kinds of models ``radiative'' and ``convective'' models in this paper. For both cases, we assume the deviation of the models from isentropic structure is small. Examining the dissipation timescales due to the gravitational radiation and several viscous processes for the polytropic neutron star models, we find that the gravitational radiation driven instability of the nodeless $r$-modes associated with $l'=|m|$ remains strong even in the non-isentropic models, where $l'$ and $m$ are the indices of the spherical harmonic function representing the angular dependence of the eigenfunction. Calculating the rotational modes of the radiative models as functions of the angular rotation frequency $\\Omega$, we find that the inertial modes are strongly modified by the buoyant force at small $\\Omega$, where the buoyant force as a dominant restoring force becomes comparable with or stronger than the Coriolis force. Because of this property we obtain the mode sequences in which the inertial modes at large $\\Omega$ are identified as $g$-modes or the $r$-modes with $l'=|m|$ at small $\\Omega$. We also note that as $\\Omega$ increases from $\\Omega=0$ the retrograde $g$-modes become retrograde inertial modes, which are unstable against the gravitational radiation reaction. ",
        "introduction": "Since the recent discovery of the gravitational radiation driven instability of the $r$-modes by Andersson (1998) and Friedman \\& Morsink (1998), a large number of papers have been published on the properties of the rotational modes of rotating stars because of their possible importance in astrophysics (e.g., Kojima 1998, Andersson, Kokkotas \\& Schutz 1998, Kokkotas \\& Stergioulas 1998, Lindblom \\& Isper 1998, Beyer \\& Kokkotas 1999, Lindblom, Mendell \\& Owen 1999, Kojima \\& Hosonuma 1999, Lockitch \\& Friedman 1999, Yoshida \\& Lee 2000, Lockitch 1999, Yoshida et al 1999). Here, we have used the term of ``rotational mode'' to refer to both $r$-mode and inertial mode of rotating stars. As an important consequence of the instability, for example, Lindblom, Owen \\& Morsink (1998) argued that due to the instability, the maximum angular rotation velocity of hot neutron stars is strongly restricted. Owen et al. (1998) also suggested that the gravitational radiation emitted from hot young neutron stars due to the $r$-mode instability would be one of the potential sources for the gravitational wave detectors, e.g. LIGO.    The dominant restoring force for the rotational modes is the Coriolis force and the characteristic frequency is comparable to the angular rotation frequency $\\Omega$ of the star (e.g., Greenspan 1964, Pedlosky 1987). The $r$-modes are characterized by the properties that the toroidal motion dominates the spheroidal motion in the velocity field, and that the oscillation frequency $\\omega$ observed in the corotating frame of the star is given by $\\omega=2m\\Omega/(l'(l'+1))$ in the limit of $\\Omega\\rightarrow0$, where $m$ and $l'$ are the indices of the spherical harmonic function $Y_{l'}^m$ representing the toroidal component of the velocity field of the mode (e.g., Papaloizou \\& Pringle 1978 , Provost, Berthomieu \\& Rocca 1981, Saio 1982). On the other hand, the inertial modes have the comparable toroidal and spheroidal components of the velocity field and therefore they are not necessarily represented by a single spherical harmonic function (e.g, Lee, Strohmayer \\& Van Horn 1992, Yoshida \\& Lee 2000). Note that the inertial mode of this kind is also called ``rotation mode'' or ``generalized $r$-mode'' (Lindblom \\& Ipser 1998, Lockitch \\& Friedman 1999).    Most of the recent studies on the $r$-mode instability driven by the gravitational radiation reaction are restricted to the case of neutron stars with isentropic structure, since the barotropic (one-parameter) equation of state, $p=p(\\rho)$ is usually assumed for both the interior structure and the small amplitude oscillation. In isentropic stars the specific entropy of the fluid in the interior is constant and the buoyant force has no effects on the oscillation. However, stars in general have non-isentropic interior structure, which may be divided into two kinds of layers with fluid stratification stable or unstable against convection. It is well know that the property of the oscillation modes of non-isentropic stars is substantially different from that of isentropic stars, since the buoyant force comes into play as a restoring force for low frequency oscillations like $g$-mode propagating in the layers with the stable stratification (see, e.g., Unno et al 1989). This is also the case for rotational modes, which have low oscillation frequencies comparable to the spin rate $\\Omega$ of the star. For example, we note that the $r$-mode associated with $l'=|m|$ having the nodeless eigenfunction in the radial direction is the only $r$-mode permitted in an isentropic star for a given $m$ (e.g, Lockitch \\& Friedman 1999, Yoshida \\& Lee 2000), but in non-isentropic stars with the stable fluid stratification all the $r$-modes with $l'\\ge |m|$ are possible for a given $m$, and there are in principle an infinite number of $r$-modes for given $l'$ and $m$ (e.g., Provost et al 1981, Saio 1982, Lee \\& Saio 1997). We therefore consider it important to investigate the modal properties of the rotational modes in non-isentropic stars, particularly in order to answer the question whether the property of the $r$-mode instability for neutron star models with non-isentropic structure remains the same as that for isentropic neutron star models.     In this paper, we study the $r$-mode and the inertial mode oscillations of slowly rotating, non-isentropic, Newtonian stars, where uniform rotation and adiabatic oscillation are assumed to simplify the mathematical treatments. The effects of the rotational deformation of the equilibrium structure on the eigenfunction and eigenfrequency are included in the analysis, using the formulation by Yoshida \\& Lee (2000) (see also Lee \\& Saio 1986, Lee 1993). We employ simple polytropic models with the polytropic index $n=1$ as the background neutron star models. For the non-isentropic models we consider only two cases, that is, the models that have stable fluid stratification in the whole interior and those that are in convective equilibrium. In this paper, following the conventions used in the field of stellar evolution and pulsation, we simply call these two kinds of models ``radiative'' and ``convective'' models, respectively, although no energy transport in the interior is considered. The plan of this paper is as follows. In \\S 2, we briefly describe the basic equations employed in the linear modal analysis of slowly rotating stars. In \\S 3, we show numerical results for the rotational modes of non-isentropic stars, and discuss the properties of the eigenfrequency and eigenfunction of the modes. In \\S 4, we examine the stability of simple neutron star models against the $r$-modes, computing their growth or damping timescales due to the gravitational radiation reaction and some viscous damping processes. \\S 5 is for conclusions. ",
        "conclusions": "In this paper, we have numerically investigated the properties of rotational modes for slowly rotating, non-isentropic, Newtonian stars, where the effects of the centrifugal force are taken into consideration to calculate the modes with accuracy up to the order of $\\bar\\Omega^3$. Estimating the dissipation timescales due to the gravitational radiation and the shear and bulk viscosity for simple polytropic neutron star models with the index $n=1$, we have shown that the gravitational radiation driven instability of the $r_0$-mode with $l'=|m|$ is hardly affected by introducing the deviation from the isentropic structure, and that the instability of the $r_0$-modes with $l'=|m|=2$ remains strongest even for the non-isentropic models. Note that the $r_0$-modes with $l'=|m|$ are the only $r$-modes possible in isentropic stars and their modal property does not very much depend on $\\Omega$ and the deviation from isentropic structure. We thus conclude that the result of the stability analysis of the $r_0$-modes with $l'=|m|$ for isentropic stars remains valid even for non-isentropic stars.   We have numerically found that for the radiative models the inertial modes at large $\\bar\\Omega$ become identical with the $r_k$-modes with $l'=|m|$ and $k>0$ and $g$-modes in the limit of $\\bar\\Omega\\rightarrow 0$ because of the effect of the buoyant force, where the integer $k$ denotes the number of node of the eigenfunction with dominant amplitude. This behavior of the $g$-modes or the inertial modes as functions of $\\Omega$ has been suggested by, for example, Saio (1999) and  Lockitch \\& Friedman (1999). We find the connection rules between the $r_k$-modes with $l'=|m|$ and $k>0$ and $g$-modes at small $\\bar\\Omega$ and the inertial modes at large $\\bar\\Omega$. We also suggest that the $r_k$-modes with $l'>|m|$ and $k\\ge 0$, whose frequency in the corotating frame becomes vanishingly small as $\\bar\\Omega$ increases, have at large $\\bar\\Omega$ no corresponding inertial modes of the isentropic models.   Good progress has been made in the understanding of the properties of the $r$-mode instability due to the gravitational radiation since its discovery. Recently, for example, Yoshida et al (1999) have shown that the results of the $r$-mode instability obtained by using the slow rotation approximation, as employed in this paper, are consistent with the results obtained by calculating directly the $r$-mode oscillations of rotating stars without the approximation. Kojima (1998) and Kojima \\& Hosonuma (1999) have studied the time evolutional behavior of $r$-modes by using the Laplace transformation technique for slowly rotating relativistic models. Lockitch (1999) also obtained $r$-modes and rotation modes for slowly rotating, incompressible, relativistic stars. However, we believe it necessary to show that the gravitational radiation driven instability of $r$-modes remains sufficiently robust in more realistic and unrestricted situations, taking account of the effects of differential rotation, magnetic field and general relativity, for example. To investigate the $r$-mode instability in more general situations is one of our future studies."
    },
    "0002/astro-ph0002136_arXiv.txt": {
        "abstract": "We constructed a model of infrared and sub-mm (hereafter IR) galaxy number count and estimated history of the IR luminosity density. We treat the evolutionary change of galaxy luminosities as a stepwise nonparametric form, in order to explore the most suitable evolutionary history which reproduces the present observational results. We found the evolutionary patterns which satisfy both constraints required from Cosmic Infrared Background (CIRB) and IR galaxy number counts. One order of magnitude increase of luminosity at redshift $z=0.75 - 1.0$ was found in IR $60\\; \\mu$m luminosity density evolution. We also found that a large number of galaxies ( $\\sim 10^7$ in the whole sky) will be detected in all-sky survey at far-infrared by {\\sl Infrared Imaging Surveyor} ({\\sl IRIS}); Japanese infrared satellite project Astro-F. ",
        "introduction": "Recent infrared and sub-mm surveys revealed a very steep slope of galaxy number count compared with that expected from no evolution model, and provided a new impetus to the related field. Such excess of galaxy number count is generally understood as a consequence of strong galaxy evolution, i.e. rapid change of the star formation rate in galaxies. Now Japanese infrared satellite project Astro-F ({\\sl Infrared Imaging Surveyor}: {\\sl IRIS}) is in progress, and we calculated the expected number count by a simple empirical method (Takeuchi et al. 1999; Hirashita et al. 1999). The applied model was based on the {\\sl IRAS} surveys, and we need more realistic one to study the detailed observational plans and follow-up strategies. In order to construct the advanced model, we first compiled the infrared/sub-mm SEDs of galaxies obtained by {\\sl ISO}, SCUBA, and other facilities, and derived average SEDs for various classes of galaxies. Then, using the local luminosity function, we studied the required galaxy evolutionary history statistically.   \\begin{figure}[ht] \\resizebox{\\hsize}{!}{\\includegraphics{ishiit1.ps}} \\caption{Assumed galaxy spectral energy distribution (SED) in the near infrared to radio wavelengths. \\label{fig1}} \\end{figure}   \\begin{figure}[ht] \\resizebox{\\hsize}{!}{\\includegraphics{ishiit2.ps}} \\caption{The double power-law form (Soifer et al. 1987) for the local luminosity function. We assumed pure luminosity evolution in this study. \\label{fig2}} \\end{figure}  \\begin{figure}[ht] \\resizebox{\\hsize}{!}{\\includegraphics{ishiit3.ps}} \\caption{Stepwise nonparametric form of the evolutionary change of galaxy luminosities. \\label{fig3}} \\end{figure}  \\begin{figure}[ht] \\resizebox{\\hsize}{!}{\\includegraphics{ishiit4.ps}} \\caption{The expected cosmic infrared background (CIRB) spectra with Evolution $1-3$ and No evolution. The observational constraints on CIRB obtained by COBE measurement (Fixsen et al. 1998, Hauser et al. 1998 and references therein) are also plotted. \\label{fig4}} \\end{figure}  \\begin{figure}[ht] \\resizebox{\\hsize}{!}{\\includegraphics{ishiit5.ps}} \\caption{{\\it The FIR $60\\; \\mu$m luminosity density evolution.} \\label{fig5}} \\end{figure}   \\begin{figure}[ht] \\resizebox{\\hsize}{!}{\\includegraphics{ishiit6.ps}} \\caption{{\\it The multiband galaxy at the infrared wavelengths.} \\label{fig6}} \\end{figure}    \\begin{figure}[ht] \\resizebox{\\hsize}{!}{\\includegraphics{ishiit7.ps}} \\caption{The multiband galaxy at the sub-mm -- radio wavelengths. \\label{fig7}} \\end{figure}  \\begin{figure}[ht] \\resizebox{\\hsize}{!}{\\includegraphics{ishiit8.ps}} \\caption{The galaxy number count predictions at assumed IRIS two bandpasses (N60 and N170). The dot-dashed lines denote the IRIS Far-infrared Scanner (FIS) flux detection limit at each waveband. \\label{fig8}} \\end{figure} ",
        "conclusions": "\\subsection{Evolutionary history}   We treat the evolutionary change of galaxy luminosities as a stepwise nonparametric form, in order to explore the most suitable evolutionary history which reproduces the present observational results. Furthermore the constraint from Cosmic Infrared Background (CIRB) should be considered as another observational constraint to the models of evolutionary history.    \\subsubsection{CIRB}  First, we searched the evolutionary pattern which satisfy the constraint required from CIRB and we found three patterns (Evolution $1 - 3$) shown in Fig. 3. These three evolutionary patterns satisfy the constraint from CIRB (Fig. 4). In order to satisfy the high background intensity at $\\sim 150\\; \\mu$m, the high evolutionary factor at $z \\sim 0.8$ is a mandatory. We note that too large evolutionary factor at high $z$ would produce the excess around $1000\\; \\mu$m.   We obtain IR $60\\; \\mu$m luminosity density along with redshift (Fig. 5) from Fig. 3. Figure 5 show the rapid evolution in $\\rho_{\\rm L} (60\\; \\mu{\\rm m})$. The increase is well described by $(1 + z)^5$~! We need such a sudden rise of $\\rho_{\\rm L} (60\\; \\mu{\\rm m})$ to reproduce the very high CIRB intensity at $\\sim 150\\; \\mu$m mentioned before. The peak of the IR luminosity density is located at $z \\sim 1$.    \\subsubsection{number count}   Recently, new observational results of the galaxy number counts at the mid-infrared and far-infrared (mainly by {\\sl ISO}), and submillimeter (SCUBA and others). We are able to compile these data as well as previously obtained {\\sl IRAS} data and radio data. It is obviously important to calculate multiband number count predictions and compare the multiband observations, because the response of the results to the galaxy evolutionary form varies with different wavelengths. In principle, the galaxy number count is an integrated value along with redshift $z$, and the information of the redshift is not available. But the redshift degeneracy can be solved to some extent, by treating the multiband observational results at the same time.  We check whether the three evolutionary patterns found in the previous section also satisfy the constraints from observations of number counts. We compare our number counts with observations in Fig. 6 (infrared), and Fig. 7 (sub-mm -- radio). Every pattern satisfys the constraints. When we especially focus on the sub-mm number counts, the evolution 3 is the most desirable.    \\subsection{Infrared Imaging Surveyor ({\\sl IRIS})}  Infrared Imaging Surveyor ({\\sl IRIS}) is a satellite which will be launched in 2003, by the M-V rocket of the {\\sl ISAS} (the Institute of Space and Astronautical Science in Japan). One of the main purposes of the {\\sl IRIS} mission is an all-sky survey at far-infrared (FIR) with a flux limit much deeper than that of {\\sl IRAS}. Detailed information of {\\sl IRIS} is available at http://koala.astro.isas.ac.jp/Astro-F/index-e.html.  In order to examine the performance of the survey, we estimated the FIR galaxy counts in two narrow bands (i.e. N60 and N170) based on models described in the previous section. We found that a large number of galaxies ( $\\sim 10^7$ in the whole sky) will be detected in this survey (Fig. 8)."
    },
    "0002/astro-ph0002408_arXiv.txt": {
        "abstract": "In this work we analyze the beryllium-iron chemical diagram from the point of view of non-LTE effects. Be abundances were re-calculated by considering non-LTE corrections in ionization equilibrium (logg) and Fe abundances ([Fe/H]). These corrections seem do not affect the linear relation between Be-Fe for metal-poor stars already found in the literature for LTE derived abundances. ",
        "introduction": "The analysis of the trends of abundances of light elements with respect to [Fe/H] for the oldest metal-poor stars is a direct way to provide some clues on their production mechanism and evolution.  In a recent work on non-LTE effects in iron abundances, Th\\'evenin \\& Idiart (1999) (TI99) obtained that for metal-poor dwarf stars Fe abundances ([Fe/H]) are affected by significant non-LTE effects and, moreover, surface gravities (logg) derived by LTE analysis also need corrections. This logg corrections should be crucial for beryllium abundances determination, since Be II resonance lines normally used to estimate Be abundances are much sensitive to this stellar parameter.  In this work we examine the consequences of non-LTE corrections to logg and [Fe/H] for logN(Be/H) vs. [Fe/H] (or Be-Fe) diagram. In section 2 we present a short summary of our results obtained in TI99 for Fe and in section 3 the results for Be abundances . ",
        "conclusions": ""
    },
    "0002/astro-ph0002314_arXiv.txt": {
        "abstract": "In 1997, the BL Lac Object Mrk 501 entered  a very active phase and was the brightest source in the sky at TeV energies, showing strong and frequent flaring. Using the data obtained with a high density air shower array that has been operating successfully at Yangbajing in Tibet since 1996, we searched for $\\gamma$-ray signals from this source during the period from February through  August in 1997. Our observation detected multi-TeV $\\gamma$-ray signals at the 3.7~$\\sigma$ level during this period. The most rapid increase  of the excess counts was observed between April 7 and June 16 and  the statistical significance of the excess counts in this period was  4.7~$\\sigma$.   Among several observations of flaring TeV $\\gamma$-rays from Mrk 501 in 1997, this is the only observation  using a conventional air shower array. We present the energy spectrum of $\\gamma$-rays which will be worthy to compare with those obtained by imaging atmospheric Cerenkov telescopes. ",
        "introduction": "Mrk 501 and Mrk 421 have been well detected as extra-galactic TeV $\\gamma$-ray sources by Whipple and subsequent ground-based Cerenkov detectors (\\cite{ong98}). They are the so-called BL Lac objects, which are radio-loud AGNs (Active Galactic Nuclei) whose relativistic jets are aligned along our line of sight. Flux variability on various scales is a common feature of BL Lac objects as already seen in Mrk 501 and Mrk 421, and spectral variations of $\\gamma$-rays coming from these sources are considered to be a very powerful tool for understanding the physics of BL Lac objects. When Mrk 501 was first detected by the Whipple Collaboration in 1995 (\\cite{quin96}), it showed rather low fluxes at a level significantly below the Crab flux. In March of 1997, however, this source went into a state  of remarkably flaring activity and its high state lasted for almost half a year with highly variable and strong $\\gamma$-ray emission. The maximum flux  reached roughly 10 times that of the Crab. During this period, several groups (\\cite{prot97}) observed strong $\\gamma$-ray emission from this source with imaging atmospheric Cerenkov detectors.  Independent measurements of the $\\gamma$-ray spectrum  seem to show a gradual softening towards higher energy, while the systematic uncertainties in the flux estimates remain too large to reach a common understanding. The energy spectrum and its shape are very important quantities for clarifying the mechanism of $\\gamma$-ray production or particle acceleration at the source, and eventually to lead to the actual measurement  of the intergalactic infrared or far-infrared background field (\\cite{steck93}). Hence,  confirmation of the detection of $\\gamma$-rays with a different technique will be strongly required.  The Tibet air shower array, operating since 1990, is located at Yangbajing (4300 m above sea level) in Tibet (\\cite{ame92}). This array has a capability of detecting $\\gamma$-rays in the TeV energy region with high efficiency and good angular resolution. Using this array, we  have succeeded in detecting the Crab at the 5.5~$\\sigma$ level (\\cite{ame99}). In this paper we  present the observation of multi-TeV  $\\gamma$-ray flares  from Mrk 501 in 1997. The result obtained with well established air shower technique is important for comparing  with those by imaging atmospheric  Cerenkov telescopes. ",
        "conclusions": "A circular window  was used to search for signals and then its size  was  determined based on the  angular resolution estimated by the experiment. The window size is chosen to optimize the significance of signals defined by $N_s/N_B^{1/2}$, where $N_S$ is the number of signals and $N_B$ the number of background events, and to contain  more than 50~\\% of the signals from a source. The radii of search windows used for the events with $\\sum\\rho >$ 15, 50 and 100 were  0.9$^\\circ$, 0.8$^\\circ$  and 0.5$^\\circ$, respectively. The signals were searched for by counting the number of events coming  from the on-source window.  The background was estimated  by averaging over events falling in the ten off-source  windows adjacent to the source, but without  overlapping each other. The source window traverses a path in local coordinates expressed by the zenith angle and azimuth angle through every day. In order to reduce a strong zenith angle dependence of the background, these off-source windows were taken in the azimuth angle directions with the same zenith angle, except two  windows adjacent to the on-source window.  \\placefigure{fig3}  Figure 3 shows the cumulative excess counts for all events  as a function of MJD and background, obtained with the HD array. No excess counts were observed  until the middle of March 1997. However, excess events  rapidly increased  in the period from  April through June and then it became slightly dull. The operation of the array was stopped on August 25 of 1997 to calibrate the operation system.  As discussed in  \\S 3, one should first note that the observed excess counts are by no means due to some artificial noise or unstable operation of the system. The statistical significance of the excess counts  reached a 3.7~$\\sigma$ level during this period. The excess counts very rapidly increased during the period from  April 7 through June 16 and  the statistical significance of the excess counts was a 4.7~$\\sigma$. These observed features are almost consistent with other observations by atmospheric Cerenkov telescopes (\\cite{prot97}).  \\placefigure{fig4}  Shown in Fig. 4 is the contour map of the excess event densities around Mrk 501 for the events with $\\sum\\rho >$ 15 observed between April 7 and June 16 in 1997. This map was obtained using the same method as done for the Moon and Sun shadows (\\cite{ame93}). Mrk 501 is well observed in the right direction by our air shower array.  \\placefigure{fig5}  Figure 5 shows the distribution of the opening angles relative to the Mrk 501 direction for all events with $\\sum\\rho >$ 15 in the HD array. The excess in the small opening angle region (less than 0.5$^\\circ$) could be attributed  to  $\\gamma$-rays from Mrk 501. The simulation result done for $\\gamma$-ray events coming from Mrk 501 can well reproduce the experiment as shown in Fig.~5, when we take account of the systematic pointing errors estimated in \\S 3. For the observation period from 1997 February to 1997 August, the statistical significances of the excess events with $\\sum\\rho >$ 15, 30 and 50 were 3.7~$\\sigma$, 2.3~$\\sigma$ and 1.6~$\\sigma$, respectively.  We also searched for $\\gamma$-ray emission using the entire Tibet-II array, but no excess was found in this period and upper limits on the excess number of the events at the 90 \\% confidence level were obtained.  We  estimated the $\\gamma$-ray spectrum from Mrk 501 by a Monte Carlo simulation (we used a GENAS code by Kasahara \\& Torii (\\cite{kasa91}).), assuming a differential power-law spectrum with the form $E^{-\\beta}$  and the cut-off at a certain energy, $E_c$, where the cut-off means that the spectral slope steepens by 1.0 at $E_c$.  The value of $\\beta$ was changed between 2.4 and 2.7 and also the effect of $E_c$ was examined  between 7 TeV  and 30 TeV. Primary $\\gamma$-rays with energies between 0.2 TeV and 50 TeV  were thrown from the direction of Mrk 501.  Observation of simulated events at Yangbajing level was done as  in our experiment, estimating the collecting area, trigger efficiency and threshold energy for $\\gamma$-rays generating the events at observation level. Simulated events in respective size ($\\sum\\rho$) bins were then compared with those by the experiment. The energy of $\\gamma$-rays was defined as the energy of the maximum flux of simulated events observed in each size bin.  These steps were repeated until the observed results are well reproduced.  A combination of $\\beta \\cong 2.6$ and $E_c \\sim$ 20-30 TeV can reproduce the data well.  We examined that the absolute flux values except the highest energy bin stay  almost unchanged for above trials, but  it is of course difficult to settle the spectral slope  from this experiment because of  very small energy range fitted here.  The  systematic errors on the flux arise mainly from the event selection procedure, which  depends upon the array performance, and from the calculations of the collecting area and the air shower size distribution by the simulation. They are estimated to be 13 \\% and 8 \\%, respectively (\\cite{ame99}).  \\placefigure{fig6}  \\placefigure{fig7}  Shown in Figs. 6 and 7 are the energy spectra averaged in the period from February 15 to August 25 in 1997 and  from February 15  to June 8 in 1997, respectively. The latter observation time corresponds to that of the Whipple Collaboration (\\cite{samu98}). It is seen that the results reported recently by other experiments (\\cite{samu98} ; \\cite{haya98} \\& \\cite{kon99}) are almost compatible with ours, although these do  not cover the same observation times. It should pay attention, however, that our results were obtained by the continuous observation of Mrk 501 extending  February through August in 1997, while those by Cerenkov telescopes were obtained for very limited  periods  of  moonless and cloudless nights.  Mrk 501 and Mrk 421, nearby AGNs, are at  almost the same red-shift (0.033 and 0.031, respectively) and have been detected  in TeV energies (\\cite{ong98}). In particular, Mrk 501 during the strong, long-lasting 1997 flare provided a good opportunity to study the energy spectrum of $\\gamma$-rays from this source in detail (\\cite{prot97}),  suggesting  a  spectral feature different with that of Mrk 421 (\\cite{kren99}). In both sources, it is likely that a synchrotron-inverse Compton picture plays an important part (\\cite{prot97}).   Since the attenuation mechanism of TeV $\\gamma$-rays by intergalactic infra-red photon field  is almost the same for both sources, a  difference of spectral features, if any,  could be attributed to the production mechanism of $\\gamma$-rays at the sources. Therefore, it is very important to continue  the observation  of high energy $\\gamma$-rays  from both sources with as small uncertainties as possible."
    },
    "0002/astro-ph0002064_arXiv.txt": {
        "abstract": "I calculate the noise in the measured correlation functions and spectra of digitized, noiselike signals. In the spectral domain, the signals are drawn from a Gaussian distribution with variance that depends on frequency. Nearly all astrophysical signals have noiselike statistics of this type, many with important spectral variations. Observation and analysis of such signals at millimeter and longer wavelengths typically involves sampling in the time domain, and digitizing the sampled signal. (Quantum-mechanical effects, not discussed here, are important at infrared and shorter wavelengths.) The digitized noise is then correlated to form a measured correlation function, which is then Fourier transformed to produce a measured spectrum. When averaged over many samples, the elements of the correlation function and of the spectrum, follow Gaussian distributions. For each element, the mean of that distribution is the deterministic part of the measurement. The standard deviation of the Gaussian is the noise. Here I calculate that noise, as a function of the parameters of digitization. The noise of the correlation function is related to the underlying spectrum, by constants that depend on the digitization parameters. Noise affects variances of elements of the correlation function and covariances between them. In the spectral domain, noise also produces variances and covariances. I show that noise is correlated between spectral channels, for digitized spectra, and calculate the correlation. These statistics of noise are important for understanding of signals sampled with very high signal-to-noise ratio, or signals with rapidly-changing levels such as pulsars. ",
        "introduction": "\\subsection{Correlation Functions and Spectra}  Electric fields from nearly all astrophysical sources are indistinguishable from Gaussian noise. Thus, nearly all of the information in such signals lies in variances of and covariances between electric fields of different polarizations, spatial locations, or frequencies. Spectra and cross-power spectra are estimates of the variance or covariance of the electric field as a function of frequency. These spectra are the Fourier transforms of auto- or cross-correlation functions. Such correlation functions are the averaged products of pairs of elements drawn from the series, for a range of time offsets or ``lags''. The elements are drawn from separate series for cross-correlation, and from the same for autocorrelation. The resulting correlation, as a function of lag, is commonly averaged over enough realizations to provide the desired signal-to-noise ratio. The correlation function is then Fourier transformed to form the desired cross-power or autocorrelation spectrum, as a function of frequency.  In practice, correlation and averaging can take place before or after the Fourier transform; this makes little difference to the result from the standpoint of this paper.  Often, data are digitized before correlation. For bandwidths that are within the capability of digital circuitry, processing is usually more accurate and economical for digital signals than for analog signals. Digitization involves sampling, or averaging over time intervals; and quantization, or describing the signal amplitude in each interval as one of a discrete set of values, rather than as a continuous variable. Quantization is an intrinsically nonlinear operation that destroys information, unlike the linear operations of sampling and Fourier transform. I find that quantization introduces effects similar to noise in the final result, as one might perhaps expect.  Usually observers wish to minimize noise, while maintaining an invertible, deterministic relationship between the mean correlation and the underlying covariance. Optimal parameters for quantization, and errors from departures from those parameters, are topics of classic work in radio astronomy (see, for example, \\citet{coo70,hag73,kul80,dad84} and references therein). Calculation of the actual noise level can be important when signal strength varies rapidly, and quantizer settings cannot remain optimal, as is sometimes the case for pulsars \\citep{jen98}; or when the distribution of the intensity of the signal must be measured accurately \\citep{gwi00}. As sensitivities of radiotelescopes improve, and as demands on the observed data increase, calculation of the noise level from quantizer parameters can be expected to become more important.  Because correlation functions and spectra are averaged over many realizations, the Central Limit Theorem implies that the resulting correlation function or spectrum has Gaussian statistics.  Thus, the statistics of the spectrum are fully described by each spectral channel's mean and variance, and covariances between channels. The mean of the spectrum is the deterministic part of the measurement; variances and covariances are the random part, or noise.  In principle, one seeks to minimize the noise, while preserving the relationship between the mean and the underlying spectrum. In \\citet{gwi04} (hereafter Paper 1), I discussed this problem for ``white'' signals, which have zero correlation except for elements of the two series with zero lag. Here I consider the more general case, where signals have arbitrary spectral character, so that covariance can depend on lag. The effect of quantization on the statistics of such ``colored'' spectra, particularly their noise, is the subject of this paper.  I calculate the noise in the quantized cross- and autocorrelation functions. The noise differs from that for correlation of continuous data in additional terms, some of them constant and others proportional to the autocorrelation function, and products of auto- and cross-correlation functions. I present this calculation through second order in correlation.  I then Fourier transform these expressions to determine the mean spectrum and its variance. The mean spectrum is simply the Fourier transform of the mean correlation function, while the noise in the spectrum is a double Fourier transform of the noise in the correlation function. I find that in the spectral domain, a gain factor, and white noise added in quadrature, approximately represent the effects of quantization in a single channel. The added white noise is commonly known as ``quantization noise''. Indeed, the gain factor and the noise are identical to those previous workers found for the deterministic, mean spectrum \\citep{coo70,jen98}. However, I also show that noise is correlated across spectral channels.  This covariance can reduce, or increase, the total noise in the spectrum, depending on the details of the quantization scheme and the details of the spectrum. For ``white'' signals without spectral variation, the more general result of this paper reduces to that found in Paper 1. In this case, the correlations or anti-correlations of noise between channels can represent an effect of the same order as quantization noise, when integrated over all channels of a spectrum.  \\subsection{Organization of this Paper}  I consider cross-correlation of two time series, $x$ and $y$, and autocorrelation of $x$. In \\S\\ref{continuous} I introduce these underlying complex time series $x_{\\ell}$ and $y_{\\ell}$ and describe their assumed statistical properties. I calculate the mean and variance of their correlation function and its Fourier transform. I show that covariance of noise in different spectral channels is zero for correlation of the continuous (un-digitized) series.  I introduce information-destroying quantization in \\S\\ref{quantized}. Under the assumption that the covariances are small (except for the zero lag of the autocorrelation function, which must be 1), I calculate the mean and the variance of the quantized correlation function for quantized data, and present analytic expressions for them. I compare the analytical results with computer simulations and find excellent agreement.  In \\S\\ref{section_quantspectrum} I find the statistics of the cross-power and power spectra. The cross-power spectrum is the Fourier transform of the cross-correlation function; and the power spectrum (sometimes called the autocorrelation spectrum) is the Fourier transform of the autocorrelation function. I calculate the noise in the spectra by a double Fourier transform of the noise in the correlation functions. I show that the noise in a single spectral channel can be approximately represented by a gain factor and white ``digitization noise'' added in quadrature with the original signal. However, I also show that noise is correlated (or, more commonly, anti-correlated) across spectral channels. I present analytical results for autocorrelation functions, and autocorrelation spectra, in \\S\\ref{acf_subsection} and\\ \\ref{acspect_subsection}. I summarize results in \\S\\ref{summary}, and show that correlations of noise between channels can represent an effect of the same order as quantization noise, when integrated over all spectral channels. ",
        "conclusions": "\\label{discussion}  \\subsection{Quantization Noise: One of Many Channels}  In the limit of spectra with many channels, $2N>>1$, the noise in one particular channel is given by terms in Eqs\\ \\ref{digital_ckrkckrk} and\\ \\ref{digital_ckrkckrkc} with coefficient $2N$ for cross-power spectra. In this approximation, \\begin{eqnarray} \\la \\ckr_k \\ckr_k^*\\ra &\\approx & \\la \\ckr_k\\ra \\la\\ckr_k^*\\ra +{{2N}\\over{2N_o}} \\big[ 2 B_X^2 B_Y^2 \\big] \\left(\\twa_k+\\left({{A_{X2}}\\over{B_X^2}}-1\\right)\\right) \\left(\\twa_k+\\left({{A_{Y2}}\\over{B_Y^2}}-1\\right)\\right)\\\\ \\la \\ckr_k \\ckr_k\\ra &\\approx & \\la \\ckr_k\\ra \\la\\ckr_k\\ra +{{2N}\\over{2N_o}} \\big[ 2 B_X^2 B_Y^2 \\big]\\twp_k\\twp_k \\end{eqnarray} These equations closely resemble the expressions for noise for continuous correlation, Eqs.\\ \\ref{continuous_spectrum_facts} and\\ \\ref{continuous_spectrum_facts}, except that everything has been multiplied by the gain factor $B_X^2 B_Y^2$, and a white noise component with variance $({{A_{X2}}\\over{B_X^2}}-1)$ (or the corresponding quantity for $Y$) has been added to the autocorrelation spectrum $\\twa_k$. These factors are those represented in the gain of the quantized cross-power spectrum (see Eq.\\ \\ref{avg_hat_rk}), and in the gain and offset of the autocorrelation spectrum (see Eq.\\ \\ref{avg_hat_rk}). Note that $A_{X2}>B_X^2$ for all $(v_0,\\;n)$, so that the added noise component is always positive. This component is conveniently interpreted as quantization noise. In this particular approximation, treatment of the effects of quantization as white noise added in quadrature is accurate.  \\subsubsection{Correlation of Noise and Noise Reduction}  Because the noise in different spectral channels is covariant (often with negative covariance), the integrated noise across a spectral channel is different from the summed, squared values of the noise in each channel (often less). Eqs.\\ \\ref{covar_pkplc},\\ \\ref{covar_pkpl}, and\\ \\ref{covar_akal} give the covariances. Although the covariances are smaller than the variances of the spectral channels given above by factors of $2 N$, they sum coherently across the channel, whereas the variances do not.  Thus, in principle they yield comparable contributions when summed over all channels. In practice, of course, the results of such a sum are given by Eqs.\\ \\ref{digital_rtrtc} and\\ \\ref{digital_rtrt} with $\\tau=0$, or Eq.\\ \\ref{a0a0c} for autocorrelation, because the sum over all spectral channels yields the zeroth lag. The interested reader can verify that the results for this lag are identical to those of Paper 1, for a white spectrum.  In principle, the reduction of noise by the covariances offers the possibility of reducing quantization noise in a spectrally-narrow signal.  For example, one could introduce additional correlated signals, with known $\\tilde \\alpha_k$ and $\\tilde \\rho_k$, and measure the variation of those from theoretically-expected results. Using Eqs.\\ \\ref{covar_pkplc},\\ \\ref{covar_pkpl}, and\\ \\ref{covar_akal} one can calculate what weighted sum of those variations should be applied to the unknown signal, to reduce the noise as much as possible. This potential application is closely related to ``dithering'' in quantization (see, for example, \\citet{bal05} and references therein).  \\subsection{Symmetries}  Note that the noise in the cross-correlation function depends on both $\\alpha$ and $\\rho$. The variance $\\la\\ckr_k\\ckr_k^*\\ra-\\la\\ckr_k\\ra\\la\\ckr_k^*\\ra$ measures the summed squares of the principal axes of the elliptical Gaussian distribution of noise, or its overall size. As one might expect from Eq.\\ \\ref{continuous_twrktwrkc}, that size depends primarily on the autocorrelation spectrum in that channel $\\twa_k$. The error ellipse must maintain the same size under the transformation $\\rho\\rightarrow e^{i\\phi}\\rho$, so the noise can depend only on even powers of $\\twa$, as it does.  Similarly, the variance $\\la\\ckr_k\\ckr_k\\ra-\\la\\ckr_k\\ra\\la\\ckr_k\\ra$ measures the difference of the squares of the principal axes of the elliptical Gaussian distribution of noise, or its shape. This shape must be circular for $\\twp=0$, so that the variance must vanish there, and so one expects that it cannot depend on $\\twa$ independently of $\\twp$. Eq.\\ \\ref{continuous_twrktwrk} confirms this. The difference must remain the same under the transformation $\\twp\\rightarrow-\\twp$, for example, so dependence on $\\twp$ must be second order.  \\subsection{Limits of Validity}  Numerical experiments suggest that Eqs.\\ \\ref{digital_ckrkckrkc} and\\ \\ref{digital_ckrkckrk} reach their limits most commonly for spectra encountered in radio astronomy when the autocorrelation function becomes large at lags other than the zero lag. For example, for a single narrow line, when the integrated power in the line becomes comparable to the integrated continuum (including system noise), then the autocorrelation function will reach about 0.5 in nonzero lags. This usually leads to noise larger than that expected from the second-order analytical expressions, especially in channels containing the line, but also throughout the spectrum.  For particular spectra, the additional noise can be modeled accurately by expressions that involve higher-order terms allowed by the preceding discussion, such as $\\twa_k\\twp_k^2$ or $\\twa_k^3$."
    },
    "0002/astro-ph0002252_arXiv.txt": {
        "abstract": "The gamma-ray pulsar PSR B1509$-$58 and its surrounding nebulae have been observed with the CANGAROO 3.8m imaging atmospheric \\v{C}erenkov telescope. The observations were performed from 1996 to 1998 in Woomera, South Australia, under different instrumental conditions with estimated threshold energies of 4.5 TeV (1996), 1.9 TeV (1997) and 2.5 TeV (1998) at zenith angles of $\\sim$ 30$^{\\circ}$. Although no strong evidence of the gamma-ray emission was found, the lowest energy threshold data of 1997 showed a marginal excess of gamma-ray--like events at the 4.1 $\\sigma$ significance level. The corresponding gamma-ray flux is calculated to be $\\rm (2.9\\,\\pm\\,0.7) \\times 10^{-12}\\,cm^{-2}s^{-1}$ above 1.9 TeV. The observations of 1996 and 1998 yielded only upper limits (99.5\\% confidence level) of $\\rm 1.9 \\times 10^{-12}\\,cm^{-2}s^{-1}$ above 4.5 TeV and $\\rm 2.0 \\times 10^{-12}\\,cm^{-2}s^{-1}$ above 2.5 TeV, respectively. Assuming that the 1997 excess is due to Very High-Energy (VHE) gamma-ray emission from the pulsar nebula, our result, when combined with the X-ray observations, leads to a value of the magnetic field strength $\\simeq$ 5 $\\mu$G. This is consistent with the equipartition value previously estimated in the X-ray nebula surrounding the pulsar. No significant periodicity at the 150\\,ms pulsar period has been found in any of the three years' data. The flux upper limits set from our observations are one order of magnitude below previously reported detections of pulsed TeV emission. ",
        "introduction": "\\label{INTRO} Pulsar nebulae have been suggested as a possible acceleration site of high-energy particles in the galaxy (\\cite{har90}). The first order Fermi acceleration mechanism is expected to occur in a shock between the pulsar wind and supernova ejecta, or interstellar matter. Evidence of such energetic phenomena has been obtained through observation of synchrotron emission by accelerated electrons and positrons at radio to gamma-ray ($\\leq$ 10 GeV) energies. However, more direct evidence has become obtainable through Very High-Energy (VHE) gamma-ray ($\\geq$ 300 GeV) observations over the last decade using the Imaging Atmospheric \\v{C}erenkov Technique (IACT).  VHE gamma-ray emissions from the directions of three energetic pulsars, the Crab (\\cite{wee89}; \\cite{vac91}; \\cite{tan94}); the Vela pulsar (\\cite{yos97}) and PSR B1706$-$44 (\\cite{kif95}, \\cite{cha97}), have been detected by ground-based telescopes using the IACT. Although all three pulsars show pulsed emission in the {\\it EGRET} energy range (100MeV--10GeV), none of the VHE gamma-ray detections have shown any periodicity at the radio pulsar period. This steady VHE gamma-ray emission is usually explained to be a result of the inverse Compton scattering in the pulsar nebula, and not from the pulsar magnetosphere. While the mechanism of the emission from the Crab nebula is well studied (see, for example, \\cite{jag96}), information on other pulsars is still sparse. In order to study pulsars and their surrounding environment as possible acceleration sites of the cosmic rays, more examples in the VHE gamma-ray range are required.  PSR B1509$-$58 was discovered as an X-ray pulsar by Seward and Harnden (1982) using the {\\it Einstein X-ray Observatory}. It is near the center of the supernova remnant MSH15$-$52 (G320.4$-$1.2). Soon after this discovery, pulsed radio emission was found by Manchester, Tuohy and D'Amico (1982). The pulsar has a period of 150 msec and a period derivative of 1.5 $\\times$10$^{-12}$ ss$^{-1}$, the largest known today. The characteristic age of the pulsar is estimated to be $\\sim$1700 years (\\cite{man98}), which makes it the second youngest pulsar after the Crab. \\altaffilmark{1} \\altaffiltext{1} { Torii et al.\\ (1997) have reported the discovery of a pulsar 1600 years old. This age is somewhat speculative however as the period derivative of the pulsar has not yet been measured and an association with a historical supernova was assumed to estimate the pulsar age.} From the period and the large period derivative, a very strong surface magnetic field of 1.5 $\\times$ 10$^{13}$ G and a large spin down energy loss rate of 1.8 $\\times$ 10 $^{37}$ ergs s$^{-1}$ are implied. While the distance to the pulsar is relatively large (4.4 kpc, \\cite{tay95}), the expected energy flux received at the Earth is the fifth largest among the known pulsars.  A compact ($\\rm\\sim\\,10^{\\prime}\\,\\times\\,6^{\\prime}$) synchrotron X-ray nebula has been found to exist around PSR B1509$-$58 (\\cite{sew84}). The synchrotron emission suggests the existence of non-thermal electrons (positrons) in the nebula, which will also emit VHE gamma-rays via inverse Compton scattering. A detectable VHE gamma-ray flux from this synchrotron nebula was predicted by du Plessis et al.\\ (1995) as a function of the magnetic field strength in the nebula. The expected gamma-ray flux above 1 TeV of 10$^{-11}$ to 10$^{-12}\\,cm^{-2}s^{-1}$ for nebula magnetic fields 4 to 10 $\\mu$G is within the sensitivity of the CANGAROO 3.8m telescope. Thus, VHE observations should give a good measurement of the magnetic field strength of this nebula. Du Plessis et al.\\ (1995) also predicted a very hard differential spectral index of $\\sim$1.8 based on the X-ray observations. This prediction provides us with an extreme example of the utility of multiwavelength studies of synchrotron---inverse- Compton emitting objects. Besides the compact nebula, recent X-ray satellite observations suggest various non-thermal phenomena in this remnant. {\\it ROSAT} observations indicate a non-thermal X-ray component from the central diffuse nebula (CDN) extending to a diameter of 50$^{\\prime}$ ($\\sim$ 60pc) centered on the pulsar (\\cite{tru96}). {\\it ASCA} observations revealed a non-thermal jet structure between the pulsar and the center of a thermal nebula about 10$^{\\prime}$ north from the pulsar (\\cite{tam96}). In order to explain the effective thermalization process of the thermal nebula, Tamura et al.\\ (1996) indicate the existence of accelerated ions as well as electrons in the jet. Furthermore, Gaensler et al.\\ (1998) found synchrotron emission from compact knots in this thermal nebula from 20cm imaging observations with the Australia Telescope Compact Array.  The surface magnetic field strength of the pulsar PSR B1509$-$58 is estimated to be one of the largest among known pulsars. Due to the photon splitting process caused by this strong surface magnetic field, a cut-off in the pulsed emission around MeV energies is predicted by Harding, Baring and Gonthier (1997). In fact, Kuiper et al.\\ (1999) have suggested that a cut-off around 10 MeV exists in the {\\it COMPTEL} data. {\\it EGRET} observations have resulted in only an upper limit for the pulsed emission from PSR B1509$-$58 (\\cite{tho94}). In contrast, Nel et al.\\ (1992) have reported the detection of transient pulsed VHE gamma-rays from the observations between 1985 and 1988 based on ground--based (non-imaging) \\v{C}erenkov telescope observations. However they could not detect any significant pulsed emission in the successive years. They tried to explain their observations with the framework of the outer gap model (\\cite{chr86}). (Bowden et al.\\ (1993) reported a upper limit of the pulsed emission above 0.35 TeV from their observations in 1987 and 1989. Combining with the detection by Nel et al.\\ (1992) in 1987 above 1.5 TeV, power law index of the integral energy spectrum is limited to be harder than $\\sim$ 1.) Interestingly, Kuiper et al.\\ (1999) also indicate a marginal detection of the pulsed emission above 10 MeV, where the origin may differ from that at lower energies. Consequently, we have examined our data for the presence of periodicity as well. Our observations are the first results on this pulsar with using the IACT, which is one order of magnitude more sensitive than non-imaging observations.  For the reasons given above, we believed that PSR B1509$-$58 would be an interesting object to study above 1 TeV energies with the CANGAROO 3.8m IACT telescope in both the steady nebula emission and the pulsed emission. Details of those observations are given in Section-\\ref{OBS}. The methods of the analysis and results are shown in Section-\\ref{ANA}. In Section-\\ref{DIS}, we summarize our results and discuss their implications. ",
        "conclusions": "\\label{DIS} Our observations can be summarized as follows $\\colon$ (1) In the observations with the lowest detection threshold energy, a 4.1$\\sigma$ excess of gamma-ray--like events is found. Null results in the observations of the other years (when the detection threshold energies were higher) are not in conflict with this marginal positive result $\\colon$ neither variability of the source nor an especially soft energy spectrum needs to be invoked. (2) From the result in the 1997 observations, there is no evidence of a variability on a monthly time-scale during three observation seasons. (3) In the 1997 data, the peak emission source position is shifted slightly to the  south-west direction from the pulsar position. However, considering the statistical error including the real event numbers observed, this is consistent with the pulsar position. (4) The periodicity of the events modulated with the radio pulsar period is studied. We found no evidence of the 150~ms pulsar periodicity using the H-test in any of the observations for three years.  The statistical significance of the 1997 excess, 4.1$\\sigma$, is too small to claim as the detection of a VHE gamma-ray source, however, it is sufficiently suggestive to allow discussion supposing the excess was due to a VHE gamma-ray signal. With this scheme the simplest and most straightforward explanation can be made assuming that the emission is found from the pulsar nebula surrounding the pulsar. VHE gamma-ray emission from a pulsar nebula is usually considered as a result of inverse Compton scattering by relativistic electrons. From the emission processes of synchrotron and inverse Compton radiations, a simple equation, $\\rm \\frac{\\dot{E}_{synch}}{\\dot{E}_{iC}}=\\frac{\\epsilon_{B}}{\\epsilon_{ph}}$, can be obtained. Here $\\rm \\dot{E}_{synch}$ and $\\rm \\dot{E}_{iC}$ are the luminosities through synchrotron radiation (mainly resulting in quanta in the X-ray energy range) and inverse Compton scattering (mainly producing VHE gamma-rays), respectively, and $\\rm\\epsilon_{B}$ and $\\rm\\epsilon_{ph}$ are the energy densities of the magnetic field and the target photons for inverse Compton scattering at the emission region. Assuming isotropic emission of both X-rays and gamma-rays, $\\rm \\frac{\\dot{E}_{synch}}{\\dot{E}_{iC}}$ can be equated to $\\rm \\frac{F_{synch}}{F_{iC}}$. Here $\\rm F_{synch}=7.2\\times10^{-11}~ergs~cm^{-2}\\,s^{-1}~(0.1-2.4\\,keV)$ as given by Trussoni et al.\\ (1996) and $\\rm F_{iC}=2.7\\times10^{-11}~ergs~cm^{-2}\\,s^{-1}$, obtained by integrating the 1997 flux above 1.9\\,TeV assuming a differential spectral index of 2.5. (The corresponding luminosity at the pulsar, $L_{iC}$, is $\\rm 6.2\\times10^{34}~ergs~s^{-1}$ assuming the pulsar distance of 4.4 kpc. That is 0.34\\% of the pulsar rotating energy loss.) If the 3\\,K Microwave Background Radiation (MBR) is the only target of the inverse Compton radiation, {\\it i.e.}, $\\rm \\epsilon_{ph}=\\epsilon_{3K}=3.8\\times10^{-13}~ergs~cm^{-3}$, one obtains $\\rm \\epsilon_{B}= 1.0\\times10^{-12}~ergs~cm^{-3}$. This, then, leads to a value for the magnetic field strength $\\rm B\\,\\simeq\\,5\\,\\mu G$. Considering the large uncertainties in the arguments above, this value agrees well with the previously estimated value of $\\rm B\\,\\simeq\\,7\\,\\mu G$, from the equipartition of energy between the particles and the magnetic field (\\cite{sew84}). According to the prediction of du Plessis et al.\\ (1995), our result corresponds to a magnetic field strength of $\\rm B~\\simeq~5~\\mu G$. These three estimated values of the magnetic field agree very well with each other.  An alternative source of the target photons is the IR source IRAS 15099$-$5856, known to be positionally coincident with the pulsar (\\cite{are91}). Du Plessis et al.\\ (1995) estimated that the contribution from the IR photons to the VHE gamma-ray flux would be at the same level as that from the 3\\,K MBR. However the association between IRAS 15099$-$5856 and the pulsar is uncertain. In case that the IRAS source found at 25$\\mu$m supplies the target photon for the inverse Compton process, the resultant VHE gamma-ray spectrum is expected to be softer than that made from the 3\\,K MBR. This is because the critical energy of the parent electrons in the Klein-Nishina cross section is $\\sim~6~\\times~10^{12}$ eV against 25$\\mu$m IR radiation while it is $\\sim~10^{15}$ eV for the 3\\,K MBR. Therefore, the VHE gamma-ray spectrum should have a rapid softening over the TeV energy range. To understand the association of this IRAS source, detailed spectral measurements with future observations are required as well as the X-ray observations discussed below.  While our observations do not place any interesting limit on the spectral index, the very hard spectrum predicted by du Plessis et al.\\ (1995) should be discussed. Their prediction was based on the observational results of the X-ray spectrum which showed a hardening of the index in the energy range below a few keV (photon index 1.4$^{+0.4}_{-0.2}$ below 4\\,keV while 2.15$\\pm$0.02 between 2 keV and 60 keV). However, recent X-ray observations do not confirm this hardening. The photon indices obtained in the wide X-ray energy band are consistent with a value around 2.2 (\\cite{tru96}; \\cite{tam97}; \\cite{mar97}) though the error of the ROSAT result is large. To discuss the synchrotron spectrum in detail, we need information from radio observations. But, even with the recent high resolution observations, a radio pulsar wind nebula has not been discovered (\\cite{gae98}).  The upper limits set to the periodic signal in this paper are one order of magnitude below the previously reported flux in the same energy band (\\cite{nel92}). Although Nel et al.\\ reported upper limits from observations after 1988, our results should provide a far stricter limit on models. The VHE pulsed emission is in conflict with the observed cut-off around 10 MeV as predicted by the polar-cap model. To explain the VHE pulsed emission, an additional hard component, probably outer-gap emission, is required. Future observations by GLAST may reveal the existence of this component and studies of its flux and spectral variability may hint at large variability in the VHE range. The flux of the transient VHE pulsed emission reported in 1985, $\\rm F(E\\,\\geq\\,1.5\\,TeV)=(3.9\\,\\pm\\,0.9) \\times 10^{-11}~cm^{-2}~s^{-1}$, would make this source the brightest known VHE gamma-ray source in the southern hemisphere. We could detect this kind of activity even with short duration monitoring. Semi-simultaneous monitoring of this pulsar with the future large IACT arrays in the southern hemisphere (CANGAROO-III, HESS) and GLAST would be of great interest if the pulsar were to display such an active phase in the future.  Finally, it is notable that, unlike the other pulsar nebulae detected at VHE energies, PSR B1509$-$58 is not firmly detected by {\\it EGRET} onboard the {\\it CGRO} satellite. In contrast, this pulsar and its surroundings show a variety of the non-thermal phenomena as introduced in Section-\\ref{INTRO}. A comparison of nonthermal X-ray emission with VHE gamma-ray emission is becoming very useful in the search for VHE gamma-ray sources and study of their environment. Combined with the recent studies of pulsar nebulae (\\cite{kaw96}), the new generation of the Imaging Atmospheric \\v{C}erenkov Telescopes (e.g.\\ \\cite{mat97}) will result in an improved understanding of pulsar nebulae and particle acceleration. The CANGAROO~II 7m telescope started observations at Woomera in mid-1999. From new observations with a lower energy threshold, we will be able to measure the gamma-ray spectrum precisely and obtain a better estimation of the physical parameters, especially the magnetic field strength, in pulsar nebulae."
    },
    "0002/astro-ph0002228_arXiv.txt": {
        "abstract": "We develop a  neutrino transport theory, in terms of the real-time non-equilibrium Green's functions, which is applicable to physical conditions arbitrary far from thermal equilibrium. We compute the coherent neutrino radiation in cores of supernovae by evaluating the two-particle-two-hole (2p-2h) polarization function with dressed propagators. The propagator dressing is carried out in the particle-particle channel to all orders in the  interaction. We show that at two loops there are two distinct sources of coherence effects in the bremsstrahlung. One is the generically off-shell intermediate state propagation, which leads to the Landau-Pomeranchuk-Migdal type suppression of radiation. We extend  previous perturbative results, obtained in the leading order in quasiparticle width,  by deriving the exact non-perturbative  expression. A new contribution due to off-shell finial/initial baryon states is treated in the leading order in the quasiparticle width. The latter contribution corresponds to processes of higher order than second order in the virial expansion in the  number of quasiparticles. At 2p-2h level, the time component of the polarization tensor for the vector transitions vanishes identically in the soft neutrino approximation. Vector current  thereby is conserved. The contraction of the neutral axial vector current with tensor interaction among the baryons leads to a  non-vanishing contribution to the bremsstrahlung rate. These rates  are  evaluated numerically for finite temperature pure neutron matter at and above the nuclear saturation density. ",
        "introduction": "Neutrino production in baryon encounters is among the fundamental processes by which compact stars lose their energy. The reactions  can be arranged, in general, according to the number of participating baryons, as the phase space  arguments play the central role in controlling their temperature and density dependence\\cite{CHIU_SALPETER,BAHCALL_WOLF,FLOWERS_ETAL,FRIMAN_MAXWELL,VOSKRESENSKY,PETHICK}. In the case of neutrino pair bremsstrahlung, the leading order process in the density virial expansion is the two-body reaction \\be\\label{2BODY_BR} B_1+B_2\\to B_1+B_2+\\nu_f+\\overline\\nu_f, \\ee where $B$ stands for a baryon, $\\nu_f$ ($\\bar\\nu_f$) for a neutrino (anti-neutrino)  of flavor $f = e, \\mu, \\tau$. Note that the subleading order process (i.e. the one in the absence of the spectator) vanishes for identical particles, as an on-shell propagating particle cannot radiate.  The matter in neutron stars is highly degenerate for temperatures typically below a MeV and  the elementary excitations are quasiparticles with well-defined energy-momentum relation.  Produced neutrinos are typically ``soft\" with energies of order of temperature. In this limit the intermediate quasiparticle propagator diverges as $1/\\omega$ and the amplitudes of the neutrino absorption, scattering, and emission turn out formally divergent as $1/\\omega^2$. The infrared behaviour of the in-medium rates, however,  is dominated by the neutrino phase space  factors, rather than the infrared divergence of the amplitudes and the rates of the bremsstrahlung and its space-like analogues remain finite. At the same time, at low temperatures, the contribution from the infrared region to the rate of the bremsstrahlung is negligible. The combined effect of the cancellation of the infrared divergence and the vanishing contribution from the low frequency region makes the quasiparticle approximation to eq. (\\ref{2BODY_BR}) applicable in cold neutron stars.  During the first several tens of seconds after a supernova explosion and core collapse the temperature of the dense nuclear matter is of the order of several tens of MeV. The neutrino bremsstrahlung is then suppressed, because the formation length of the neutrino radiation is of the same order of magnitude as the mean free path of a baryon\\cite{RAFFELT_SECKEL,JANKA,RAFFELT,HANESTAD}. The collective effects become important on the radiation scale (i.e. the role of the spectator in the reaction (\\ref{2BODY_BR}) is taken over by the medium) because the baryon undergoes multiple scattering during the radiation. The underlying mechanism is the Landau-Pomeranchuk-Migdal (LPM) quenching of the radiation, first  introduced in the context of QED\\cite{LP}. The central role in the theory is played by the formation length of radiation $l_{\\rm f}$. If the mean-free-path of a baryon is much larger that the formation length $l_{\\rm mfp}\\gg l_{\\rm f}$ then the radiation  reduces to a sum of separate radiation events, each of which is well described by the  Bethe-Heitler spectrum. In the opposite limit $l_{\\rm mfp}\\ll l_{\\rm f}$ the individual scattering events are unresolved and the radiation spectrum takes the Bethe-Heitler form for a  single scattering event. In the intermediate regime, when $l_{\\rm mfp}\\sim l_{\\rm f}$, the radiation amplitudes for scattering off various centers interfere destructively and the radiation is suppressed (Landau-Pomeranchuk-Migdal effect; for a review see refs. \\cite{RAFFELT,KNOLL_VOSKRESENSKY}).  The rates of neutrino-nucleon processes are commonly expressed through phase space   integrals over the contraction of the weak currents with  the polarization function of the nuclear medium. The polarization function (or structure function) of the supernova/neutron star matter has been subject of many studies\\cite{RAFFELT,IWAMOTO_PETHICK,SAWYER,HOROWITZ,HAENSEL,BURROWS,REDDY_ETAL}. The modifications  of reactions rates by the spatial correlations among (on-shell) quasiparticles have been studied within the Fermi-liquid theory\\cite{IWAMOTO_PETHICK}, the one-boson exchange interaction theory\\cite{SAWYER}, the relativistic random phase approximation\\cite{HOROWITZ},  the variational approach\\cite{HAENSEL}, and combinations thereof\\cite{BURROWS,REDDY_ETAL}. The spatial correlations tend to suppress reaction rates in general, although their impact on the supernova physics is model dependent\\cite{IWAMOTO_PETHICK,SAWYER,HOROWITZ,HAENSEL,BURROWS,REDDY_ETAL}.   The common strategy of incorporating the LPM effect in the neutrino-nucleon interaction processes is to add a quasiparticle damping in the intermediate state propagator by replacing $\\omega$ by $\\omega +i\\gamma$\\cite{RAFFELT_SECKEL,JANKA,RAFFELT}. In the soft neutrino limit  the vector current coupling does not contribute by virtue of the vector current conservation (CVC) and the net  contribution comes from the axial-vector transitions via baryon spin-flip. The above modification of the intermediate state propagator then leads to an  ansatz for the nucleon spin structure function: $S_{\\sigma}\\propto \\gamma_{\\sigma}/(\\omega^2 + \\gamma_{\\sigma}^2)$ \\cite{RAFFELT_SECKEL,JANKA,RAFFELT},  where $\\gamma_{\\sigma}$ is the nucleon spin-flip collision rate. The ansatz generalizes the  quasiparticle picture, in a semi-phenomenological manner,  by including the temporal correlations among the quasiparticles in the leading order in the quasiparticle width.  The microscopic justification of this phenomenology emerges from the various formulations of the finite temperature quantum filed theory, e. g. the thermo-field dynamics\\cite{BRAATEN} or the closed diagram formalism in the Schwinger-Keldysh technique\\cite{KNOLL_VOSKRESENSKY}. A microscopic computation is not straightforward, however. For example, the polarization function of the medium can be computed at one-loop, including the quasiparticle width to all orders in $\\gamma$, however {\\it a priori} the current conservation is not guaranteed at this level.  The reason, in part, is that the  ``more complicated'' higher order in  loop expansion diagrams contribute  at the same order as the single loop\\cite{KNOLL_VOSKRESENSKY}.  In a previous paper we carried out a microscopic computation of the bremsstrahlung, including the LPM effect, at the one-loop level in a formalism based on the quasiclassical Kadanoff-Baym transport equation\\cite{SD}. Here we extend this computation to two-loops and partially modify our approach to include the propagator and vertex renormalization on the same footing and  to including the tensor force explicitly. The extension to two-loops is motivated by the following. The long range phenomena, driven by the weaker attractive part of the baryon-baryon interaction, are sensitive to the resummation in the particle-hole ($ph$) channel. On the other hand, as well known,  one should fully resum the particle-particle ($pp$) channel to treat the hard core of the baryon-baryon interaction. Therefore, the $ph$ channel can be treated perturbatively by expanding in the number of particle-hole loops, while the $pp$ channel must be treated non-perturbatively by a full resummation of the ladder diagrams. Thus, the separation of the long range and short range phenomena dictates the manner in which the diagrammatic expansion is carried out. The dressing of the single particle propagators occurs in both channels and can be treated either explicitly, say, by considering higher order self-energies attached to a propagator, or, alternatively, by condensing it in the width of the propagator spectral function. As a consequence of the separation of the scales, the short-range correlations can be condensed in the propagator width on the scales relevant for the long-range phenomena. The imaginary part of a single loop in the $ph$ channel vanishes in the time-like region of the phase space, which is relevant for the particle production. A  finite result emerges when one dresses the  propagators by either extending the resummation in the $ph$ channel to two and higher loops and/or by dressing the propagators in the $pp$ channel to all orders. Ignoring the latter resummation, i.e. using the quasiparticle propagators in the two-loop expansion, misses a number of short-range collective effects, such as the LPM quenching of the radiation due to multiple scattering. On the other hand, summing only the ladders in the $pp$ channel does not recover  the vector current conservation in the radiation process (in, at least,  a transparent manner). Therefore a natural choice, motivated by the separation of the short and long range phenomena, is to truncate the $ph$ channel at two-loops and to resum the $pp$ channel to all orders. The situation is reminiscent of the  parquet resummation scheme in the first iteration, where in both channels the driving force is the bare baryon-baryon interaction.   Early studies of the bremsstrahlung at the quasiparticle level modelled the strong force using the $T$-matrix interaction\\cite{FLOWERS_ETAL}, the free-space one-boson exchange interaction\\cite{FRIMAN_MAXWELL} and their in-medium modifications \\cite{VOSKRESENSKY} supplemented with a hard core modelled in the spirit of the Fermi-liquid theory. The explicit use of the tensor interaction turned out to be crucial as there are significant cancellations among different diagrams, and the surviving contribution is due to a non-trivial contraction between the operator structures of the weak and strong interactions (tensor force)~\\cite{FRIMAN_MAXWELL}. This motivates our ansatz for the driving force in the particle-hole ($ph$) channel of nuclear interaction, which includes explicitly the tensor force contribution. We do not attempt, in the present work, to go beyond the one-pion exchange approximation for several reasons, one being that the non-perturbative treatment of the interaction does not change the spin, isospin, and tensor operator structure of the interaction, and important cancellations in the radiation matrix elements will be preserved in a more advanced treatment. We also want to be able to isolate the finite width effects in our comparisons to the earlier work done in the one-pion exchange approximation\\cite{FRIMAN_MAXWELL,RAFFELT}. The situation is different in the particle-particle ($pp$) channel, where the short-range correlations have to be treated in a non-perturbative manner by summing up the ladder diagrams to all orders. We do this in the finite-temperature Brueckner theory.   The paper is organized as follows. In Section 2, starting from the Kadanoff-Baym formalism, we derive a single-time transport equation for (anti)-neutrinos with collision integrals driven by (anti)-neutrino coupling to baryons via the polarization tensor of the medium. The polarization tensor is computed in the 2p-2h approximation in Section~3. The summation of the ladder diagrams in the $pp$ channel within the finite temperature Brueckner theory is described in Section 4. Section~5 evaluates the phase space integrals and neutrino bremsstrahlung emissivities. The numerical results are presented in Section~6. Section 7 summarizes our main results. ",
        "conclusions": "In this work we formulated a transport theory for neutrinos in the framework of real-time Green's functions formalism, with particular attention to the collision integrals for the neutrino-pair bremsstrahlung. The main focus was a first principle  calculation of the bremsstrahlung emissivity including the width of propagators. This allows to answer questions, not covered by the semi-phenomenological theory, such as the magnitude of the contribution of higher order terms in the expansion with respect to the quasiparticle width or the cancellation of the  vector current contribution at all orders in the quasiparticle width. Even though the expression for the emissivity, which follows from our quasiclassical transport equation, is the same as the one found in the linear response theory, it is valid under conditions arbitrary far from equilibrium. This is particularly important in the regime where the  neutrinos decouple from matter and their distribution function strongly deviates from the Fermi-Dirac form.   The central quantity of the theory is the particle-hole polarization tensor in the $ph$ channel truncated at two loops. The $pp$ channel is treated non-perturbatively within the finite temperature Brueckner theory. We find that the only contribution to the bremsstrahlung rate comes from the contraction of the tensor force with the axial vector current to all orders in the quasiparticle width. Other contributions, which arise from the contraction of the  Fermi-liquid type interaction with the axial vector current and the contraction of the net strong interaction with the vector current, cancel when we take the sum of the diagrams corresponding to vertex corrections and propagator renormalization in the $ph$ channel. Thereby the vector current conservation is established at all orders in the quasiparticle width. These cancellations are independent of the approximations to the propagators and are effective both in the quasiparticle limit and beyond. The three ingredients crucial to the cancellations are: (i) the anti-commutation of the tensor force with the axial vector current, (ii) the odd parity of the causal propagator under the exchange of its energy argument, (iii) the soft neutrino and non-relativistic baryon approximations.    Our numerical evaluation of the neutrino emissivity of hot neutron matter, carried out at two loops, shows that the LPM-type suppression sets in at temperatures $T\\ge \\gamma$, in agreement with the previous work limited to the first order terms in the quasiparticle width (see ref. \\cite{RAFFELT} and references therein). The higher order terms enhance the magnitude of the neutrino emissivity compared to the leading order result. The non-perturbative result, however, is still suppressed as compared to the quasiparticle limit.  Our formalism can be extended in various ways. One obvious extension is allowing for two different chemical potentials of scattering baryons. This will include the Urca process (the $\\beta$-decay in the second order in the virial expansion) and the effects of the Pauli spin-paramagnetism, which become important in strong magnetic fields. The formalism can be adapted, with minor changes, for a computation of the space-like analogous of the bremsstrahlung and, in particular, the neutrino opacities of the supernova matter. The perturbative scheme, employed here,  itself requires further improvements in several  direction, numerically the most  important one being the renormalization of the one-boson exchange interaction in the $ph$ channel."
    },
    "0002/astro-ph0002472_arXiv.txt": {
        "abstract": " ",
        "introduction": "Gamma-ray bursts have puzzled astronomers since their accidental discovery in the late sixties. Recent observations still provide new topics of study, afterglow,\\cite{Cosra} optical flash,\\cite{Akerlof} brightest events,\\cite{Castro} and so forth. In these bursts, huge amounts of energy are released instantaneously. The  subsequent evolution of the radiation-dominated matter can be described by an expanding blast wave. Gamma-rays are emitted when a matter with highly relativistic velocity, with Lorentz factor $\\Gamma \\geq 10^2 $, is converted into radiation.  The observed fluctuations in gamma-rays with a short timescale are thought to be associated with the complex behavior of shock waves. Two locations have been proposed as the origin of the irregularities. One location is inside the relativistic fluid, as described by the internal shock model. When shock waves intersect each other, gamma-rays are emitted. The complexity of the corresponding time profile  depends on  the internal shock structure. The other location is in the shock front itself, as described by the external shock model. When the shock front impinges upon ISM (inter stellar matter), gamma-rays are emitted. Both models have advantages and disadvantages. Waxman\\cite{waxman-a}\\cite{waxman-b} showed synchrotron emission from external shocks provides a successful model for the broken power law spectra and smooth temporal behavior of afterglow. Sari and Piran,\\cite{Sari1} however, showed that external shocks cannot produce variable gamma-ray bursts unless they are produced by extremely narrow jets or if only a small fraction of the shell emits radiation. The expanding shell and ISM were assumed to be smooth. The observed nature with rapid variability depends on the fluctuations of the external matter, which is not yet known. The possibility of irregularity in the shock front can be investigated. An irregular shell with angular fluctuations in the external shock may produce temporal behavior of the gamma-ray emission.  A theoretical model of the blast wave with ultra-relativistic speed was considered by Blandford and McKee.\\cite{Blandford} They obtained a spherical  self-similar solution, which can be regarded as the relativistic version of the Sedov-Taylor solution. Using a numerical code for the spherical symmetry, Kobayashi, Piran and Sari\\cite{Sari2} simulated an expanding shock with relativistic speed and confirmed that the self-similar solution is a good approximation. The spherical symmetry may be too great a simplification in some cases. Indeed, in non-relativistic hydrodynamics, Ryu and Vishniac\\cite{Ryu} showed that under certain conditions the spherical shock becomes overstable with respect to perturbations with tangential velocity. From this point of view, it is important to examine the dynamical stability of the Blandford-McKee solution.  In this paper, we consider the linear stability with respect to non-spherical deformations of ultra-relativistic shock waves. There is a similarity solution describing shocks with an impulsive energy supply. In Section 2,  we show that the perturbation functions can be obtained in a power low form. Matching the boundary conditions at the shock front and using a regularity condition at the origin, the temporal behavior, i.e. an eigenvalue of the functions, is determined. Section 3 is devoted to discussion. ",
        "conclusions": "In this paper, we have studied the linear stability of a spherically expanding shock in the ultra-relativistic limit, i.e. the large $ \\Gamma $ limit. The temporal behavior of the perturbation is expressed by a power law in $ \\Gamma$. We found a solution that grows with the deceleration of the shock. The growth rate becomes large for short wavelengths of the angular fluctuation. That is, there exists a rapidly growing mode for the non-spherically symmetric displacement. Vishniac\\cite{Vishniac} discussed the physical interpretation of the unstable mode in the non-relativistic blast wave. It is important to remember that two kinds of pressures exerted to the shock plane have completely different natures. In the exterior region before the shock, the ram pressure dominates and is parallel to the direction of the shock propagation. After the shock, the thermal pressure dominates and is isotropic. When the shock spherically expands, two directions, i.e. the direction of the propagation and the normal direction of the plane coincide. However, when the shock is rippled with the tangential velocity, two pressures lose the balance in the tangential direction. This unbalance is expected to be larger with the increase of the velocity. The unstable mechanism indeed works in the ultra relativistic blast wave as shown in this paper. This result depends on several assumptions and approximations, but it is quite suggestive. Suppose, e.g., for the relativistic flow dynamics from $ \\Gamma = 10^2$ to 1 associated with the gamma-ray bursts, that the amplitude increases by more than a factor of $10^{20}$. The dominant scale of the unstable mode depends on both the initial disturbance and the growth rate. If there is no characteristic wavelength in the initial stage of the expansion, the unstable mode is likely to be non-linear at the smaller scale due to the large growth rate. This mode may lead to a number of fragmentations. The radiation from the blobs may be reflected by an irregular time-profile of the gamma-rays.  In order to check this possibility, further study is necessary. For dynamical evolution, we must simulate relativistic flows using axi-symmetric or 3-dimensional numerical codes. Fluctuations are manifest only in the small region after the shock front. The structure may be found only by using a fine resolution in the numerical calculations. The unstable mode may be suppressed to a certain value of the amplitude. Also, the size and number relations of the fragmentation can be estimated. The effect of the flow dynamics will appear in the emission, but the  gamma-ray emission process is not clear at moment. However, the complex structure of the shock front may universally give the time history of the burst, irrespective of the detailed behavior of the emission process.  \\appendix"
    },
    "0002/astro-ph0002191_arXiv.txt": {
        "abstract": " ",
        "introduction": "Attempts to understand star formation in low surface brightness (LSB) galaxies has resulted in a large number of theories being discarded and few alternatives being offered.  As a result we have considerable knowledge on what these enigmatic systems are not.  LSB galaxies are {\\em not}:  \\begin{itemize}  \\item simply the faded version of high surface brightness (HSB) galaxies.  Although some red LSB galaxies have been found which may be the end product of the faint blue galaxies, the majority of LSB galaxies have very blue colors and low metallicities (i.e.  Ferguson \\& McGaugh 1995; O'Neil, \\etal 1997a; McGaugh 1994; Schombert, \\etal 1990; De Blok \\& Van der Hulst 1998), arguing against any fading scenario.  \\item lacking the neutral hydrogen necessary to form stars, as many LSB galaxies contain more than 10$^9$ \\Msol of H I and LSB galaxies include some of the highest M$_{HI}$/L$_B$ galaxies known (O'Neil, Bothun, \\& Schombert 1999).  \\item a completely new type of galaxy.  The transition from HSB to LSB galaxies is smooth, with LSB galaxies covering the entire color and morphological spectrum of HSB galaxies (i.e.  O'Neil, \\etal 1997b; Matthews \\& Gallagher 1997)  \\end{itemize}  UGC 12695 is a relatively nearby (z=0.021) low surface brightness galaxy with an absolute blue magnitude of M$_B$=$-$18.9.  Previous studies of UGC 12695 (McGaugh, 1994; O'Neil \\etal, 1998) have shown it to be very remarkable.  The galaxy is of an exceedingly transparent nature, evidenced by the many background galaxies seen through its elusive disk, and it contains a reasonably high gas fraction (M$_{HI}$/L$_B$ = 2.6 \\Msol/\\Lsol) while having a very low metallicity and an extremely blue color for a galaxy ($U-I$ = $-0.2$) (Table 1).  Because UGC 12695 was thought to be fairly isolated, with the nearest galaxy (UGC 12687) lying more than 277 kpc away (Figure~\\ref{fig:Allposs2HI}), it provides a good opportunity for studying star formation and evolution in LSB galaxies.  To this end, and with the above points in mind, we undertook to observe UGC 12695 with the Very Large Array (VLA) in the C configuration.  The results of these observations are described in this paper, as follows: Section 2 describes the observations and data reduction; Section 3 examines the H I morphology and kinematics of UGC 12695 and its companions --  UGC 12687, and 2333+1234; Section 4 looks at the dark and visible mass of UGC 12695; Section 5 examines the  star formation potential of UGC 12695; Finally, section 6 examines the possibility of a recent tidal encounter between the UGC galaxies.  \\begin{table}[ht] \\begin{center} \\caption{Global properties of UGC~12695 and UGC~12687.} \\begin{tabular}{lccc} \\noalign{\\vspace{0.8mm}} \\hline \\hline \\noalign{\\vspace{0.8mm}} &            &     UGC~12695     &    UGC~12687         \\\\ \\hline \\noalign{\\vspace{0.8mm}} Type          &            &            Sm$^1$ &              SBbc$^1$\\\\ V$_{hel}$           & km s$^{-1}$&           6186$^2$&               6150$^2$\\\\ M$_B$               & mag        &          -18.9$^3$&              -20.3$^1$\\\\ $B-V$               &            &          0.26$^3$ &               0.70$^4$\\\\ $\\mu_B(0)$     &mag/arcsec$^{2}$&          23.8$^3$ &                 -     \\\\ r$_{25}$            & kpc        &          0.79$^1$ &               0.87$^1$\\\\ M$_{HI}$               & M$_\\odot$  & 7.5$\\times$10$^9$ $^2$& 1.2$\\times$10$^{10}$ $^2$\\\\ M$_{HI}$/L$_B$ & M$_\\odot$/L$_\\odot$ &      2.62$^2$ &              1.18$^2$ \\\\ \\noalign{\\vspace{0.8mm}} \\hline \\multicolumn{4}{l}{$^1$De Vaucouleurs, \\etal 1991}\\\\ \\multicolumn{4}{l}{$^2$This paper}\\\\ \\multicolumn{4}{l}{$^3$O'Neil, \\etal 1998}\\\\ \\multicolumn{4}{l}{$^4$Prugniel \\& Heraudeau 1998}\\\\ \\noalign{\\vspace{0.8mm}} \\hline \\hline \\end{tabular} \\end{center} \\end{table}   \\begin{table}[ht] \\caption{VLA observing parameters} \\begin{tabular}{lr} \\noalign{\\vspace{0.8mm}} \\hline \\hline \\noalign{\\vspace{0.8mm}} Configuration                                               &     C-short \\\\ Correlator mode                                             &  2AC-normal \\\\ Total integration time \\hfill                       (hours) &        15.5 \\\\ Dates of observation                                        &     15Jan99 \\\\ \\multicolumn{2}{r}{16Jan99} \\\\ Field center,           $\\alpha$(B1950)                     &    23:33:30 \\\\ \\phantom{Field center, }$\\delta$(B1950)                     &    12:35:53 \\\\ Central frequency \\hfill                              (MHz) &     1391.64 \\\\ $V_{\\rm hel}$ of central channel \\hfill    (km$\\:$s$^{-1}$) &        6170 \\\\ Primary beam FWHM \\hfill                           (arcmin) &        32.4 \\\\ Synthesized beam ($\\alpha$$\\times$$\\delta$) \\hfill (arcsec) & 16.2$\\times$14.1 \\\\ Bandwidth \\hfill                                      (MHz) &      1.5625 \\\\ Number of channels                                          &         256 \\\\ Channel separation \\hfill                  (km$\\:$s$^{-1}$) &        1.31 \\\\ Velocity resolution \\hfill                 (km$\\:$s$^{-1}$) &        1.58 \\\\ rms noise in one channel \\hfill                         (K) &        2.63 \\\\ K-mJy conversion,                                           &             \\\\ \\phantom{K-}equiv. of 1mJy/beam \\hfill                  (K) &        2.76 \\\\ \\noalign{\\vspace{0.8mm}} \\hline \\hline \\end{tabular} \\end{table} ",
        "conclusions": "UGC 12695 is an intriguing low surface brightness galaxy of a very transparent nature, having an extremely blue color, a highly asymmetric appearance and very localized bursts of star formation near the peaks in the H I column density distribution.  Many of the properties of both UGC 12687 and UGC 12695 can be explained as being induced by such a tidal interaction, including the bar of UGC 12687 and its central radio continuum emission and UV excess as well as the lopsided appearance of UGC 12695 and the offset between its morphological and kinematic major axes.  Furthermore, the localized bursts of star formation in UGC 12695 could very well be induced by such an interaction, giving rise to local instabilities in the LSB disk as demonstrated by Mihos \\etal (1997).  It is likely that UGC 12695 could have been living a fairly quiescent existence, its low surface gas density keeping its star formation rate quite low, and just now it is experiencing a period of localized but vigorous star formation triggered by a mild tidal interaction which might lead to a major morphological transition.  Within all this, though, it is easy to overlook one important fact. Although many of the properties of UGC 12695 and UGC 12687 can readily be explained through an ongoing tidal encounter, the two galaxies are still fundamentally distinct.  UGC 12695 is not simply a fainter, or more `stretched-out', or more quickly rotating version of UGC 12687. Were any of these the case the behavior of the two galaxies after the tidal encounter would be similar, and UGC 12695 would have experienced a central gas inflow with the majority of its star formation now occurring not in the outlying regions (as it is), but in the galaxy's core.  Thus the fundamental question of why UGC 12695 is an LSB galaxy, and UGC 12687 is not remains unanswered."
    },
    "0002/astro-ph0002158_arXiv.txt": {
        "abstract": "{  We here investigate the possibility that the ultra-high energy  cosmic ray (UHECR) events observed above the GZK limit are mostly protons accelerated in reconnection sites just above the magnetosphere of newborn millisecond pulsars that are originated by accretion induced collapse (AIC). }   \\nonstopmode  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002081_arXiv.txt": {
        "abstract": "Observations in the submillimetre waveband have recently revealed a new population of luminous, sub-mm sources.  These are proposed to lie at high redshift and to be optically faint due to their high intrinsic dust obscuration.   The presence of dust has been previously invoked in optical galaxy count models which assume $\\tau=9$ Gyr Bruzual \\& Charlot evolution for spirals and these fit the count data well from U to K. We now show that by using either a 1/$\\lambda$ or Calzetti absorption law for the dust and re-distributing the evolved spiral galaxy UV radiation into the far infra-red(FIR), these models can account for all of the `faint'($\\leq1$mJy) $850\\mu$m galaxy counts, but fail to fit 'bright'($\\ge2$mJy) sources, indicating that another explanation for the sub-mm counts may apply at brighter fluxes(e.g. QSOs, ULIRGs). We find that the main contribution to the faint, sub-mm number counts is in the redshift range $0.5 < z < 3$, peaking at $z\\approx 1.8$.  The above model, using either dust law, can also explain a significant proportion of the extra-galactic background at $850\\mu$m as well as producing a reasonable fit to the bright $60\\mu m$ IRAS counts. ",
        "introduction": "The SCUBA camera \\cite{holland99} on the James Clerk Maxwell Telescope has transformed our knowledge of dusty galaxies in the distant Universe as a result of the discovery of a new population of luminous, dusty, infra-red galaxies (Smail et al. 1997; Ivison et al 1998).  It has been proposed that these galaxies may be similar to IRAS ULIRGs (ultra-luminous infra-red galaxies) which appear to be starbursting/AGN galaxies, containing large amounts of dust.  The possibility that much star-formation is hidden by dust means that sub-mm observations can give an invaluable insight into the star-formation history of the Universe.  This view aided by the redshifting of the thermal dust  emission peak in starbursting galaxies into the FIR, which results in a negative k-correction in the sub-mm. By this route, we can therefore study our Universe all the way back to very early times and gain unprecedented insight into the formation and evolution of galaxies.   The first sub-mm galaxy to be detected by SCUBA was SMM J02399-0136 \\cite{ivison98}, which is a massive starburst/AGN at z=2.8 and the current situation is that the complete 850$\\mu$m sample from all the various groups consists of well over 50  sources (Blain et al. 1999; Eales et al. 1999; Hughes et al. 1998; Holland et al. 1998; Barger et al. 1998; Smail et al. 1997).  Optical and near infra-red(NIR) counterparts have been identified for about a third of the sources, although the reliability of these identifications varies greatly.  This problem is due to the fact that the $\\approx 15''$ FWHM of the SCUBA beam results in $_-^+3''$ positional errors on a sub-mm source, so there is a reasonable chance that several candidates could lie within these errors.  Also, there is no guarantee that the true source will be detected down to the optical flux limit as, for example, many of the sources have been shown to be very red objects (Dey et al. 1999; Smail et al. 1999: Ivison et al. 2000) and therefore have not been found in optical searches for sub-mm sources.  What has proved extremely enlightening is that radio counterparts at 1.4GHz have now been identified for many of the sub-mm sources (Smail et al. 2000: Ivison et al. 2000) providing much more accurate angular positions ($<1''$ in some cases) and reasonably accurate photometric redshifts. Various groups have obtained redshift distributions of sub-mm samples (Hughes et al. 1998; Barger et al. 1999a; Lilly et al. 1999; Smail et al. 2000) and they all derive results that are consistent with a mean redshift in the range $1 < z < 3$.  The fact that almost all of the sources are associated with mergers or interactions seems to confirm that the population of sources contributing at the `bright' ($>2$mJy) sub-mm fluxes (since most of the sources so far discovered are `bright') are similar to local IRAS ULIRG's, ie massive, starbursting/AGN galaxies which are extremely luminous in the far-infra-red. This hypothesis is strengthened further by the fact that the only two sub-mm sources (SMM J02399-0136 and SMM J14011+0252) with reliable redshifts have been followed up with millimeter wave observations (Frayer et al. 1998, 1999), resulting in CO emission being detected at the redshifts of both sources (z=2.8 and z=2.6), a characteristic indicator of large quantities of molecular gas present in IRAS galaxies.  The nature of the fainter ($\\le1$mJy) sub-mm population is, however, the focus of this paper. It has been claimed by Peacock et al. (2000) and Adelberger et al. (2000) that the Lyman Break Galaxy(LBG) population could not only contribute significantly to the faint sub-mm number counts, but could also account for a substantial proportion of the background at $850\\mu$m. This may indicate that ULIRG's cannot explain all of the sub-mm population and that the UV-selected galaxy population, which are predicted to be evolved spirals by the Bruzual \\& Charlot models, may in fact make a substantial contribution. It is exactly this hypothesis our paper addresses.  In this paper we will first review the situation regarding the optical galaxy counts, focusing in particular on the models of Metcalfe et al. (1996). These simple models which use a $\\tau =9$Gyr SFR for spirals and include the effects of dust give good fits to galaxy counts and colours from U to K.  The idea is then to see whether this combination of exponential SFR and relatively small amounts of dust in the first instance (A$_B=0.3$ mag. for the $1/\\lambda$ law), which would re-radiate the spiral ultra-violet (UV) radiation into the FIR, could cause a significant contribution to the sub-mm galaxy number counts and background at $850\\mu $m.  Our modelling will be described in section 3 and then in section 4 our predicted contribution to the $850\\mu $m and $60\\mu m$ galaxy counts and the extra-galactic background in the sub-mm will be shown.  Also in this section we demonstrate how to get a fit to the background in the $100-300\\mu m$ range by using warmer, optically-thicker dust in line with that typically seen in ULIRG's. We will then discuss the implications of our predictions in section 5 and conclude in section 6. ",
        "conclusions": "The aim of this paper was to investigate whether, by re-radiating the absorbed spiral galaxy UV flux into the FIR, the dust invoked in the faint blue spirals at high z from the optical galaxy count models of Metcalfe et al.(1996) could have a significant contribution to the sub-mm galaxy counts and also the FIR background at $850\\mu$m.  We have found that, using a interstellar dust temperature of $15K$, a modest circumstellar component of $45K$, a beta parameter of 2.0 and a galaxy formation redshift of $z_f\\approx4$ we can account for a very significant fraction of the faint $850\\mu m$ source counts, both in the low and high $q_0$ cases when we invoke Bruzual \\& Charlot evolution (see Fig. \\ref{fig:counts}). These evolutionary models give 5-10 times more contribution to the faint sub-mm counts than the corresponding no-evolution models. At brighter fluxes, we find that the SFR and dust assumed in our  normal spiral model are too low  to produce the FIR fluxes of the most luminous sources.   In the no-evolution cases, we underpredict the number counts, even at the faint end.  Our predicted redshift distribution of sub-mm selected faint blue galaxies suggests that the main contribution to the faint counts is in the range $0.5 < z < 3$, peaking at $z\\approx1.8$.  We have shown that our model fits the $60\\mu m$ IRAS data well, an important local test if we want to assume PLE and extrapolate our optical spiral galaxy luminosity functions out to higher redshift. With the evolution models we can easily account for 50-100\\% of the FIR background at $850\\mu m$ but fail the data by nearly an order of magnitude in the $100-300\\mu m$ range. We have shown that the only way to fit these observations using this optically based model is to use assume more dust obscuration ($A_B=0.6$) and much warmer dust (T=30K).  Effectively gray extinction laws such as that of Calzetti et al (1997) may also provide more overall absorption and hence allow more dust temperature components to allow the flexibility to fit  the FIR background from 60-850 $\\mu m$. However, the bright  sub-mm counts will  still require a further contribution from QSO's or ULIRGs to complement the contribution of the faint blue galaxies at fainter fluxes."
    },
    "0002/astro-ph0002048_arXiv.txt": {
        "abstract": "The calculations in Thomas-Fermi approximation show that in a gravitational field each cell of ultra dense matter inside celestial bodies obtains a very small positive electric charge. A celestial body is electrically neutral as a whole, because the negative electric charge exists on its surface. On the order of magnitude the positive volume charge is very small ($10^{-18}e$ only). But it is sufficient to explain the occurrence of magnetic fields of celestial bodies and the existence of a discrete spectrum of steady-state values of masses of  stars and pulsars. ",
        "introduction": "We cannot measure magnetic fields of the majority of stars, which are distant far from us. Therefore, the existence of magnetic fields for the majority of stars can be considered only hypothetically. However, the magnetic field of the Sun is known over than a hundred years, and in the last decades the astronomers managed to measure magnetic fields for a number of stars (so-called $A_{p}$-stars) \\cite{1} and some pulsars \\cite{2}. It is interesting to construct a model describing the generation of magnetic fields by stars and to compare it with the data of the astronomers. The mechanism examined below  is based on the gravity-induced electric polarization of matter. It is capable to explain also the generation of magnetic fields by planets \\cite{3}, however, in the case of stars, this mechanism works in the purest manner. ",
        "conclusions": ""
    },
    "0002/astro-ph0002424_arXiv.txt": {
        "abstract": "Data from the Canadian Galactic Plane Survey reveal an abundance of \\hi shells and arcs in the disk of our galaxy. While their shape is suggestive of stellar winds or supernovae influence, very few of these structures have been examined in detail thus far. A fine example is an \\hi shell in the outer Galaxy with no continuum counterpart discovered in the survey's pilot project. Its size and kinematics suggest that it was created by the winds of a single late-type O star which has since evolved off the main sequence or by a supernova explosion. A B1 Ia star at the centre of the shell, in projection, is a possible candidate for energy source if the shell is assumed to be wind-blown. The shell's shape implies a surprisingly small scale height of less than about 30 pc for the surrounding gas if the elongation is due to evolution in a density gradient. ",
        "introduction": "Hot, massive stars have a major impact on the surrounding interstellar medium (ISM), not only at the end of their life when they become supernovae but also throughout their more stable phases of evolution via their strong stellar wind. The winds of O-type stars can inject as much energy into the ISM throughout their main sequence life as in their final explosion, and therefore should have an equally great impact on the structures in the ISM and its energization.  Typically, wind-blown shells have been found by looking at the environments of stars known to have strong stellar winds, e.g.\\ O stars and Wolf-Rayet stars (e.g.\\ Benaglia \\& Cappa 1999, Marston 1997, Miller \\& Chu 1993, Dubner et al.\\ 1990), but a neutral shell may continue to exist after the central star has evolved off the main sequence and has lost the power to ionize the shell. Supernova remnants (SNRs) are thought to be visible in the radio continuum for only a few to several tens of thousands of years (Braun et al.\\ 1989, Frail et al.\\ 1994). As this is considerably less than the duration of their life before they merge with the ambient ISM ($> 1$ Myr), there should be many SNRs consisting of neutral gas shells. Large scale, low resolution surveys have revealed the presence of \\hi supershells (e.g.\\ Heiles 1984), most likely created by stellar clusters and associations, but could not bring to light the smaller but likely numerous neutral shells created by single stars which are an important part of the Galaxy's zoology.  The Canadian Galactic Plane Survey (CGPS; Higgs 1999 and Taylor 1999) offers the first opportunity to study a large collection of \\hi shells, as opposed to supershells, to determine their dynamics and how they relate to and impact on the surrounding ISM. A few such objects were serendipitously discovered in the pilot project (Normandeau et al.\\ 1997; hereafter NTD97). One of these, \\g132, will be presented in detail here as an illustrative case study of this class of objects. It is a striking feature within the \\hi data cube, developing over several spectral channels, at velocities generally associated with interarm gas.  The following section briefly outlines the observations and processing of the data. Section 3 provides a description of the structure at several wavelengths. In \\S4, the possible location of the shell is discussed. The next section considers the stars present in this vicinity as possible energy sources for a wind-blown bubble. In \\S6, all these elements are brought together for analysis and conjecture. A summary and conclusions are given in \\S7. ",
        "conclusions": "An \\hi shell has been found near the very active W3/4 \\hii region complex. The lack of a radio continuum counterpart has been interpreted as indicative of the advanced age of the shell, be it a wind-blown shell or a SNR. If it is a wind-blown shell then the most likely powering source is the B1 supergiant BD+60 447. This is based on the position of the star (at the centre in projection and at a reasonable distance) and its spectral type (strong enough stellar winds while on the main sequence to blow a bubble, no longer capable of maintaining the ionization of the shell which is in accord with the lack of a continuum counterpart). The age of the star and the kinematic age of the shell are however discrepant, the former being greater than the latter. This could point to the shell being a member of the observed class of ``high velocity'' shells (Oey 1996), which have somehow been reaccelerated, but it could also be taken to indicate that the shell was not created by BD+60 447 (and there are no other catalogued stars present which could have blown the shell through stellar winds) but is in fact a SNR. Based on the available data, it is not possible to distinguish between the two possibilities.  Regardless of the origin of the \\g132, if the elongation of the shell is due to the density gradient of its surroundings, the Kompaneets model can be used to determine the scale height of the ambient ISM. In this case, the aspect ratio of the shell indicates that the local scale height is 17.3 pc for a distance of 2.2 kpc. While surprisingly low, analysis of other regions have also pointed to small scale heights (W4 by Basu et al.\\ 1999; W44 by Shelton et al.\\ 1999).  At the edge of the shell, there is a smaller U-shaped \\hi structure curving around a compact radio continuum and infrared source. The thermal spectral index of the compact source and its infrared colours imply that it is an \\hii region and the \\hi then corresponds to an encircling photodissociation region. However, the FCRAO outer Galaxy Survey shows no indication of a coincident molecular cloud at similar velocities. It has been hypothesized that this \\hii region could be the result of star formation triggered by the expansion of \\g132. Compact $^{12}$CO clouds at other locations along the shell's perimeter could also be triggered or enhanced condensations.  The data from the CGPS are likely to be rife with such structures as their identification requires arcminute resolution coupled with coverage of wide angular scales. With an analysis of this sort carried out for each shell, our picture of the star-ISM feedback mechanisms will be more complete."
    },
    "0002/astro-ph0002338_arXiv.txt": {
        "abstract": "We present observations of \\rd300, a quasar at $z = 5.50$ discovered from deep, multi-color, ground-based observations covering 74 arcmin$^2$.  This is the most distant quasar or AGN currently known. The object was targeted as an $R$-band dropout, with $R_{\\rm AB} > 26.3$ (3$\\sigma$ limit in a 3\\arcsec\\ diameter region), $I_{\\rm AB} = 23.8$, and $z_{\\rm AB} = 23.4$.  The Keck/LRIS spectrum shows broad \\lya/\\nv\\ emission and sharp absorption decrements from the highly-redshifted hydrogen forests.  The fractional continuum depression due to the \\lya\\ forest is $D_A = 0.90$.  \\rd300\\ is the least luminous, high-redshift quasar known ($M_B \\approx -22.7$). ",
        "introduction": "The past few years have witnessed a watershed in our direct observations of the high-redshift Universe.  A decade ago, only a handful of galaxies were identified past a redshift of 3.  These sources represented the rare beast in the cosmos:  high-redshift radio galaxies, or galaxies associated with extremely distant, luminous quasars.  New techniques and instruments allow us now to routinely identify normal, star-forming galaxies at these same epochs \\markcite{Steidel:96a}(\\eg Steidel {et~al.} 1996).  \\markcite{Stern:99e}Stern \\& Spinrad (1999) review modern search techniques for distant galaxies.  Improved computing power and ambitious, large-area surveys also have pushed the frontier of distant quasar studies \\markcite{Djorgovski:99, Fan:99}(\\eg Djorgovski {et~al.} 1999; Fan {et~al.} 1999).  Now we are regularly identifying objects which have collapsed only $\\sim 1$ Gyr after the Big Bang.  Such observations tell us about the earliest phases of galaxy and structure formation and probe the conditions of the early Universe.  The identification of high-redshift quasars is especially important for several reasons.  First, quasars at early cosmic epoch require the rapid formation of a supermassive black hole.  Assuming black holes are not primordial, this requires the condensation of a large cloud of hydrogen, presumably embedded within a dark matter halo.  Additionally, the presence of metal lines in quasars demand a previous generation of stars (two generations for nitrogen).  High-redshift quasars thus constrain models of galaxy and structure formation \\markcite{Loeb:93, Eisenstein:95}(\\eg Loeb 1993; Eisenstein \\& Loeb 1995).  Also, quasars provide valuable probes of the intervening intergalactic medium \\markcite{Rauch:98}(\\eg Rauch 1998) and the intergalactic ionizing background.  For example, the absence of a smooth depression in quasar continua short-ward of the \\lya\\ emission strongly constrains the amount of neutral hydrogen in the intergalactic medium \\markcite{Gunn:65}(Gunn \\& Peterson 1965).  \\markcite{Songaila:99}Songaila {et~al.} (1999) find no Gunn-Peterson trough out to redshift 5 from deep spectroscopic observations of SDSSp~J033829.31+002156.3 at $z = 5.00$ \\markcite{Fan:99}(Fan {et~al.} 1999).  In this Letter, we report the discovery of a quasar at $z = 5.50$, the most distant quasar identified to date.  The previous  most distant quasar was SDSSp~J120441.73$-$002149.6 at $z = 5.03$ \\markcite{Fan:00}(Fan {et~al.} 2000). At $z = 5.50$, an $H_0 = 50 \\kmsMpc, \\Lambda = 0, \\Omega = 1$ (0.1) universe is 790 Myr (1.51 Gyr) old, corresponding to a look-back time of 94.0\\%\\ (90.9\\%) of the age of the universe.  For the lambda cosmology supported by recent studies of distant supernovae, $H_0 = 65 \\kmsMpc, \\Lambda = 0.7, \\Omega_{\\rm m} = 0.3$, the universe is 1.11 Gyr old at $z=5.50$, corresponding to a look-back time of 92.4\\%\\ of the age of the universe. ",
        "conclusions": "Though of moderate signal-to-noise ratio, the spectrum of \\rd300\\ has the unambiguous signature of an extremely distant quasar.  The broad emission with a sharp absorption at 7900 \\AA\\ is consistent with \\lya/\\nv\\ emission attenuated by the nearly opaque \\lya\\ forest at $z = 5.50$.  An additional discontinuity is visible at 6690 \\AA, associated with the \\lyb\\ forest.  The \\lya\\ forest absorption and poor detection of the long-wavelength \\siivoiv\\ emission make centroiding on the emission features ill-advised; the redshift is instead determined from the sharp forest decrements.  We estimate $z = 5.50 \\pm 0.02$.  This and other properties of \\rd300\\ are given in Table~1.  The spectral character of \\rd300\\ is slightly atypical of high-redshift quasars, though it undoubtedly resides within the diverse category of quasars.  The \\lya/\\ion{N}{5} complex is unusually broad and distinguishing the emission lines is impractical.  Many of the highest redshift quasars share similar spectroscopic shapes, \\eg\\ SDSSp~J033829.31+002156.3 at $z = 5.00$, SDSSp~J021102.72$-$000910.3 at $z = 4.90$ \\markcite{Fan:99}(Fan {et~al.} 1999), and GB~1428+4217 at $z = 4.72$ \\markcite{Hook:98}(Hook \\& McMahon 1998).  The strong continuum absorption associated with the \\lya\\ forest is the dominant spectroscopic feature of \\rd300.  A robust determination of $D_A$, the standard parameter for describing the \\lya\\ forest decrement \\markcite{Oke:82}(Oke \\& Korycansky 1982), requires knowledge of the continuum spectral slope long-ward of \\lya.  We estimate $D_A$ by assuming the standard quasar power law spectral index of $-0.5$ for the continuum long-ward of \\lya\\ \\markcite{Richstone:80, Schneider:92}(\\eg Richstone \\& Schmidt 1980; Schneider {et~al.} 1992), with the amplitude determined over the wavelength interval $\\lambda \\lambda 8400 - 9000$ \\AA\\ (between the \\nv\\ and \\siivoiv\\ emission complexes).  We derive $D_A = 0.90 \\pm 0.02$.  We derive $D_B = 0.95 \\pm 0.04$ for the strength of the Ly$\\beta$ forest.  The $D_A$ value is comparable to those measured for distant galaxies in the Hubble Deep Field at similar redshifts \\markcite{Weymann:98, Spinrad:98}(Weymann {et~al.} 1998; Spinrad {et~al.} 1998) and models of the \\lya\\ forest \\markcite{Madau:95, Zhang:97}(Madau 1995; Zhang {et~al.} 1997).  At $z = 5.50$, the features used to describe continuum properties are redshifted to challenging wavelengths.  We estimate $AB_{\\rm 1450(1 + z)}$ using the continuum modeled above.  Consistent with previous work in this field, $M_B$ is calculated for an Einstein-de~Sitter universe with $H_0 = 50 \\kmsMpc, q_0 = 0.5$ and the standard quasar power law index of $-0.5$.  We find $M_B \\approx -22.7$; see \\markcite{Stern:00a}Stern {et~al.} (2000a) for details of how we calculate $M_B$.  Comparison with the 1.4~GHz FIRST survey \\markcite{Becker:95}(Becker, White, \\& Helfand 1995) reveals no radio source within 30\\arcsec\\ of the quasar to a limiting flux density of $f_{\\rm 1.4 GHz} \\simeq 1$ mJy (5$\\sigma$).  How unusual is it to find a quasar as distant and luminous as \\rd300\\ in a $\\simeq 100$ arcmin$^2$ field?  This is difficult to answer as \\rd300\\ is the least luminous $z > 4$ quasar known.  The previously known high-redshift ($z > 4$) quasars of lowest luminosity are PC~0027+0521 \\markcite{Schneider:94}($z = 4.21, M_B = -24.0$; Schneider, Schmidt, \\& Gunn 1994), which was discovered serendipitously, and the X-ray selected quasar RX~J105225.9+571905 \\markcite{Schneider:98}($z = 4.45, M_B = -23.9$; Schneider {et~al.} 1998). These luminosities are comparable to the lower luminosity objects in the $z \\simlt 1$ Bright Quasar Survey \\markcite{Schmidt:83}(Schmidt \\& Green 1983).  Most high-redshift quasar luminosity functions \\markcite{Schmidt:95b}(\\eg Schmidt, Schneider, \\& Gunn 1995) have been derived from samples of quasars $\\approx 100$ times as luminous as \\rd300.  To calculate the expected surface density of high-redshift, faint quasars, we follow the methodology outlined in \\markcite{Kennefick:95}Kennefick {et~al.} (1995) and \\markcite{Boyle:98}Boyle \\& Terlevich (1998):  we adopt the \\markcite{Boyle:91}Boyle {et~al.} (1991) $z = 2$ quasar luminosity function (for $q_0 = 0.5$), scaled down in density using the evolution predicted by \\markcite{Schmidt:95b}Schmidt {et~al.} (1995), namely, that the quasar space density falls off by a factor of 2.7 per unit redshift beyond $z = 3$.  The predicted surface density of $R$-drop ($4.3 \\simlt z \\simlt 5.8$) quasars with $M_B \\leq -22.5$ is $\\simeq 2 \\times 10^{-3}$ arcmin$^{-2}$, implying $\\approx 0.15$ such quasars should have been uncovered in our survey.  It is dangerous, yet enticing, to draw conclusions from a single object.  The discovery of the quasar \\rd300\\ at $z = 5.50$ in the modest sky coverage of our survey is suggestive of less dramatic evolution in the quasar luminosity function at faint magnitudes and high redshift.  Such a change would have significant cosmological implications, including changing the budget of high-energy, ionizing photons in the early Universe.  We also note that the low signal-to-noise ratio data are suggestive of strong hydrogen absorption near the quasar redshift.  Is this due to a neutral hydrogen cloud near the quasar, at odds with the proximity effect?  Or are we seeing the first glimpses of an object radiating prior to the reionization epoch, with neutral intergalactic hydrogen absorbing the rest-frame UV photons?  Higher resolution, higher signal-to-noise ratio data will be essential for answering these questions."
    },
    "0002/astro-ph0002497_arXiv.txt": {
        "abstract": "We limit the contribution of the hot intracluster gas, by means of the Sunyaev-Zel'dovich effect, to the temperature anisotropies measured by the Tenerife experiment. The data is cross-correlated with maps generated from the ACO cluster catalogue, the ROSAT PSPC catalogue of clusters of galaxies, a catalogue of superclusters and the HEAO 1 A-1 map of X-ray sources. There is no evidence of contamination by such sources at an rms level of $\\sim 8\\mu$K at 99\\% confidence level at $5^o$ angular resolution. We place an upper limit on the mean Comptonization parameter of $ y \\le 1.5\\times 10^{-6}$ at the same level of confidence. These limits are slightly more restrictive than those previously found by a similar analysis on the COBE/DMR data and indicate that most of the signal measured by Tenerife is cosmological. ",
        "introduction": "The discovery of Cosmic Microwave Background (CMB) temperature anisotropies (Smoot et al. 1992; Wright et al. 1992) immediately prompted the question of their origin. In most cosmological scenarios temperature anisotropies are expected to arise with the growth of matter density perturbations (Peebles 1980); but it was also suggested (Hogan 1992) that the anisotropies could have been originated by inverse Compton scattering from hot diffuse clouds of electrons in nearby superclusters (see Birkinshaw 1999 for a review on the Sunyaev-Zel'dovich effect). Right after Hogan's suggestion, Boughn \\& Jahoda (1993) and Bennett et al. (1993) searched for non cosmological signal contributions to the COBE/DMR 2 year sky maps. If the Sunyaev-Zel'dovich (SZ) component is significant at some level and the distribution of hot electrons is traced by rich clusters or any other extragalactic source, then temperature anisotropies should be correlated with maps constructed from extragalactic source surveys. Banday et al. (1996) performed this analysis for the COBE/DMR  4 year map. Their study did not find a statistically significant contribution, reassuring the idea that temperature anisotropies were of cosmological origin. Similar cross-correlation analysis have been carried out by several authors to estimate the rms level of the Galactic emission present in different CMB maps (Kogut et al. 1996; Kneissl et al. 1997; de Oliveira-Costa et al. 1997, 1998, 1999).  In this article we  study the contribution of nearby structures on the temperature anisotropies measured by the Tenerife map (Guti\\'errez et al. 2000). Two effects associated with the local matter distribution can give rise to anisotropies in the CMB: the non-linear evolution of gravitational structures (Rees \\& Sciama 1968) and the mentioned SZ effect. The latter is expected to give the largest contribution, and we shall center our analysis on it. Since clusters are known to posses ionized gas at temperatures about $10^8$K they are expected to give the largest contribution to the SZ effect. Cocoons of radio galaxies (Yamada, Sugiyama \\& Silk 1999) or other foreground sources subtend a very small angle and, due to beam dilution, have a very small contribution at the scales probed by the Tenerife experiment. Therefore, we shall look for correlations between the Tenerife CMB data and template maps generated from cluster catalogues; in particular, we use the ACO (Abell, Corwin \\& Olowin 1989), and ROSAT PSPC catalogues (Vikhlinin et al. 1998). To test the hypothesis that the CMB signal is generated by hot diffuse gas distributed on supercluster scales we also include a supercluster catalogue (Einasto et al. 1994) in our analysis. Finally, we also constructed template maps from the HEAO 1 A-1 map of X-ray sources (Kowalski et al. 1984) that should trace the distribution of hot electrons. In Sec. 2 we describe the statistical methods used for comparing data on temperature anisotropies with template maps constructed from different surveys. In Sec. 3 we describe the catalogues used and how the template maps were elaborated. Finally, in Sec. 4 we present and discuss our main results. ",
        "conclusions": "Table 1 summarizes the results of our analysis. After substracting the mean and normalizing the templates to unit variance, the autocorrelations and cross correlations were computed given equal weight to each pixel. We tried different angular bins, from 1$^o$ to 5$^o$, and computed the correlation function out to $20^o$ and $30^o$. No significant differences were found. In Table 1, the results are quoted for correlation functions in bins of $3^o$ out to $21^o$.  The SZ effect will generate approximately equal and negative contributions at 10 and 15GHz. If the signal detected is real one should expect equal and positive values of $\\alpha$ at the two frequencies. In Table 1 we give $\\hat\\alpha$, $\\sigma_{\\hat\\alpha}$ as given by eqs.~(\\ref{alpha}) and (\\ref{salpha}), and $\\chi^2$ per degree of freedom (dof) in the two approaches described in Sec. 2. For simplicity, we termed (a) ``without sample variance'' and (b) ``with sample variance''. In the case (a), we also give the error associated with sampling variance $\\sigma_s$, which should be added in quadrature with $\\sigma_{\\hat\\alpha}$. As the maps were normalized to unit variance, $\\alpha\\sigma_{TEN}$ gives the SZ component of the CMB map in thermodynamic units. For each template we calculate the cross-correlation with the 15GHz and 10GHz maps. Since the latter covers a smaller fraction of the sky, we also correlate a reduced 15GHz map (denoted by 15c in the table) cut to the size of the 10GHz map to eliminate the bias introduced by the different sky coverage.  No significative detections (larger than $2\\sigma$) were found by either of the two methods. We always found negative values of $\\alpha$, i.e., cold spots in the template map correlate with hot spots in the data, contrary to what one would expect if there was a significant SZ contribution. The largest signal was obtained at the 15c ROSAT template map. Since the amplitude of the SZ effect does not change much at the Tenerife frequencies, consistency would require a fluctuation of the same order to be present at 10GHz. Furthermore, when sample variance was not included in the covariance matrix, the best-fit was never a good fit. Only when it was included $\\chi^2/dof$ became of order unity. The low quality of the fit can be understood by looking at Figure 1, where we plot the autocorrelation of the template maps (Fig 1a) and their cross correlation with the data on 15GHz (Fig 1b). The dashed line corresponds to the autocorrelation and cross correlation of the ACO catalogue, the long dashed line to the HEAO 1 A-1 catalogue, the dot-dashed line to the ROSAT catalogue. Thick and thin solid lines correspond to superclusters with uniform and concentrated gas, respectively. While the autocorrelation functions are rather similar in all cases, the cross-correlation differ substantially in shape.  When the sample variance is not included in the covariance matrix, $\\sigma_{\\hat\\alpha}$ is small and the difference in shape implies a large $\\chi^2/dof$. Only when sample variances are included, the error bars are much larger and the fit to the data improves. To conclude, we can only set upper limits on the value of $\\alpha$. Taking the results on 15GHz, we limit $\\alpha\\le 0.24$ at the 99\\% confidence level.  For an experiment with such a beam width like Tenerife, one could not expect to find a large correlation between data and templates. For example, some clusters in the ACO catalogue have been found to produce temperature fluctuations of the order $100\\mu$K (Birkinshaw 1999). But they subtend an angular scale of few arcminutes and as a result the SZ signal is diluted by the large solid angle covered by the Tenerife beam. Still, the Tenerife data limits the contribution of nearby clusters and superclusters to be smaller than $8\\mu$K at 99\\% confidence level. The mean Comptonization parameter at $y = {-\\Delta T\\over 2T_o} \\le -1.5\\times 10^{-6}$ at the same level of confidence. Our results are slightly more restrictive than those previously found by Banday et al. (1996). Let us remark that $y$ obtained above only limits the contribution due to nearby superclusters, while the COBE result of  Mather et al. (1994) applies to the contribution of all structures located between the last scattering surface and the observer. To conclude, this study, like previous ones based on the COBE/DMR data, indicate that most of the signal measured by Tenerife is not of extragalactic origin but cosmological."
    },
    "0002/astro-ph0002174_arXiv.txt": {
        "abstract": " ",
        "introduction": "Powerful radio galaxies are frequently associated with extended emission line nebulosities which extend on radial scales of 5 --- 100 kpc from the nuclei of the host early-type galaxies (Hansen \\etal\\ 1987, Baum \\etal\\ 1988). The morphological and kinematical properties of these nebulae provide important clues to the origins of the gas, and the origins of the activity as a whole. The study of such gas is, for example, important for our understanding of the building of galaxy disks and bulges by infall since their epoch of formation. Therefore, it is crucial to determine the extent to which the observed emission line properties reflect the intrinsic distribution of the warm gas in the haloes of the host galaxies, and the extent to which they reflect the effects of the nuclear activity and interactions with the extended radio sources.  In most low redshift ($z < 0.2$) radio galaxies the optical emission line regions are broadly distributed in angle around the nuclei of the host galaxies, the correlations between the optical and radio structural axes are weak, and the gas kinematics are often quiescent, with line widths and velocity shifts consistent in most cases with gravitational motions in the host early-type galaxies (Tadhunter, Fosbury \\& Quinn 1989; Baum, Heckman \\& van Breugel 1990). However, optical observations reveal a dramatic change in the properties of the nebulosities as the redshift and radio power increase: the emission line kinematics become more disturbed (compare Tadhunter \\etal\\ 1989 with McCarthy, Baum \\& Spinrad 1996) and the optical/UV structures become more closely aligned with the radio axes of the host galaxies (McCarthy \\etal\\ 1987, McCarthy \\& van Breugel 1989).  The most recent high resolution HST images of $z \\sim 1$ radio galaxies show that the structures are not only closely aligned with the radio axes, but they are also highly collimated, with a jet-like appearance (Best et  al. 1996). The nature of  the ``alignment effect'' is a key issue for our general understanding of powerful radio galaxies, of particular relevance to the use of radio sources as probes of the high redshift universe.  Of the many models which have been proposed to explain the alignment effect, the two which have received the most attention are the {\\it anisotropic illumination} and the {\\it jet-cloud interaction}\\footnote{We use ``jet-cloud interaction'' as a generic term to describe interactions between the warm ISM and the radio-emitting components, which could include the radio lobes and hot-spots, as well as the radio jets.} models.  In the case of anisotropic illumination it is proposed that the gaseous haloes of the host galaxies are illuminated by the broad radiation cones of the quasars hidden in the cores of the galaxies (\\eg\\ Barthel 1989), with the emission lines resulting from photoionization of the ambient ISM by the EUV radiation in the cones (\\eg\\ Fosbury 1989), and the extended optical/UV continuum comprising a combination of the nebular continuum emitted by the warm emission line clouds (Dickson \\etal\\ 1995) and scattered quasar light (Tadhunter \\etal\\ 1988, Fabian 1989). The alignment of the obscuring tori perpendicular to the collimation axes of the plasma jets then leads to a natural alignment of the extended nebulosities with the radio axes. The best evidence to support the anisotropic illumination model is provided by polarimetric observations of powerful radio galaxies at all redshifts which show evidence for high UV polarization and scattered quasar features (\\eg\\ Tadhunter \\etal\\, 1992, Young \\etal\\ 1996, Dey \\& Spinrad 1996, Cimatti \\etal\\ 1996, Cimatti \\etal\\ 1997, Ogle \\etal\\ 1997). It would be difficult to explain these polarimetric results in terms of any mechanism other than scattering of the anisotropic radiation field of an illuminating quasar or AGN. However, despite the success of the illumination model at explaining the polarization properties, a sigificant fraction of radio galaxies --- comprising $\\sim$30 -- 50\\% of radio galaxies at $z\\sim1$, and a smaller proportion of lower redshifts --- are dominated by jet-like UV emission line structures, which are  more highly collimated than would be expected on the basis of the 45 -- 60$^{\\circ}$ opening half-angle illumination cones predicted by the unified schemes for powerful radio galaxies (Barthel 1989, Lawrence 1991). Moreover, the highly disturbed emission line kinematics observed in many high-z sources are also difficult to reconcile with quasar illumination of the undisturbed ambient ISM of the host galaxies.   Jet-cloud interactions have the potential to explain many of the features of powerful radio galaxies which cannot be explained in terms anisotropic quasar illumination. Although the jet-cloud interactions are likely to be complex, at the very least the clouds will be compressed, ionized and accelerated as they enter the shocks driven through the ISM by the jets. Therefore, jet-cloud interactions provide a promising explanation for the high-surface-brightness and extreme emission line kinematics of the structures aligned along the radio axes of high-z sources. Indeed, recent spectroscopic observations of jet-cloud interactions in low redshift radio galaxies provide clear observational evidence for the acceleration and ionization of warm clouds by the jet-induced shocks (\\eg\\ Clark \\etal\\ 1998, Villar-Martin \\etal\\ 1999). Moreover, theoretical modelling work has demonstrated that jet-induced shocks are a viable, if not unique, mechanism for producing the emission line spectra of radio galaxies (\\eg\\ Dopita \\& Sutherland 1998).  It is clear that no single mechanism can explain the emission line properties of radio galaxies over the full range of redshift and radio power; some combination of AGN illumination and jet-cloud interactions is required, with the jet-cloud interactions becoming increasingly important as the redshift and/or radio power increases. However, a major problem with such a combined model is that, while the polarimetric results demonstrate that quasar illumination is important in many high redshift sources, the extended structures are often dominated by highly collimated jet-like structures, with no sign of the  broad cone-like emission line distributions predicted by the unified schemes for powerful radio galaxies.  How do we explain this dearth of broad cones in the objects with the most highly collimated structures? Possibilities include the following. \\begin{itemize} \\item {\\bf The gas structures are intrinsically aligned along the radio axes of the high redshift sources}, so that the emission line nebulae reflect the true distribution of warm/cool gas, rather than the ionization patterns induced by the jets or illuminating AGN. For example, West (1994) has proposed that a general alignment of the gas structures along the radio axis may be a natural consequence of the formation of giant elliptical galaxies in a heirarchical galaxy formation scenario, although it is not clear that the structures formed in this way would be quite as highly collimated as those observed in the high-z radio galaxies. \\item {\\bf The nuclei of the host galaxies do not contain powerful quasars}, and the ionization of the extended gas is dominated by the jets: either by direct jet/cloud interactions, or by illumination by the relativistically-beamed jet radiation. This scenario is supported by the discovery at low redshifts of a class of powerful radio galaxies with weak, low ionization nuclear emission line regions (see Laing 1994, Tadhunter \\etal\\ 1998). However, at least some high-z radio galaxies with highly collimated UV structures show direct evidence for powerful quasar nuclei in the form of scattered quasar features in their polarized spectra, so this explanation cannot hold in every case. \\item {\\bf The broad-cone radiation of the buried quasars does not escape from the nuclear regions of the host galaxies}, because of the absorbing effects of circum-nuclear gas. In this case the ionization of the extended gas in the aligned structures is likely to be dominated by jet-cloud interactions, but quasar or beamed jet radiation may also contribute along the jet axis, if the jets punch a hole in the obscuring material. Direct evidence for requisite obscuring material in the central regions is provided by the relatively high occurrence rate of associated CIV and Ly$\\alpha$ absorption line systems in the UV spectra of radio-loud quasars (\\eg\\ Anderson \\etal\\ 1987, Wills \\etal\\ 1995), the relatively red SEDs of steep spectrum radio quasars (Baker 1997), and the detection of significant BLR reddening in a substantial fraction of nearby broad-line radio galaxies (\\eg\\ Osterbrock, Koski \\& Miller 1995, Hill \\etal\\, 1996). \\item {\\bf The dearth of broad ionization cones in the high-z sources is a consequence of an observational selection effect:} most of the existing emission line images of the high-z sources have been taken in the light of the low ionization [OII]$\\lambda$3727 line which is emitted particularly strongly by the jet-induced shocks (\\eg\\ Clark \\etal\\ 1997, 1998), but is relatively weak in  the more highly ionized quasar illumination cones. Thus, given that the published ground-based images of the high-z objects are relatively shallow and have a low spatial resolution, while the published HST images have a higher resolution but are insensitive to low surface brightness structures, the existing images are likely to be biased in favour of the high-surface-brightness shocked structures along the radio axes. In this case, we would expect deep emission images to reveal gaseous structures outside the main high surface brightness structures aligned along the radio axes. If the gas away from the radio axis is predominantly photoionized by quasars hidden in the cores of the galaxies we would expect the extended low surface brightness structures to have a broad distribution, consistent with quasar illumination. Detection of such emission line morphologies in the objects with the most highly collimated structures would lead to a reconciliation between the anisotropic illumination and jet-cloud interaction models, thereby resolving the outstanding uncertainties surrounding the nature of the alignment effect. \\end{itemize} In order to test the latter possibility it is important to obtain deep emission line imaging observations for the objects with closely aligned radio and optical structures. We report here pilot observations of two intermediate-redshift radio galaxies --- 3C171 ($z=0.2381$) and 3C277.3($z=0.08579$) --- which are nearby prototypes of the high-z radio galaxies, in the sense that they show high surface brightness emission line structures which are closely aligned along their radio axes. The results challenge some commonly-held assumptions about the ionization of the extended gaseous haloes around powerful radio galaxies. ",
        "conclusions": "The main aim of the deep emission line imaging observations was to attempt to detect the broad emission line cones outside the main aligned structures, and thereby reconcile the AGN illumination and jet-cloud interactions models. The unified schemes predict illumination cones with opening half-angles of 45-60$^{\\circ}$. Although the extended emission line nebulosities in low redshift radio galaxies rarely show the  sharp-edged cone structures observed in some Seyfert galaxies (\\eg\\ Pogge 1988, Tadhunter \\& Tsvetanov 1989), the emission line distributions are generally consistent with broad cone illumination of an inhomogeneous ISM (Hansen \\etal\\ 1987, Baum \\etal\\ 1989, Fosbury 1989). The detection of similar emission line distributions in 3C171 and Coma A would support the idea that the extended ionized haloes are photoionized by quasars hidden in the cores of the galaxies.  The deep imaging observations presented in this paper have confounded our expectations in the sense that, while they do show extended of emission line gas well away from the radio axis, the emission line distribution cannot be reconciled with any plausible ionization cone model. Not only do some of the features wrap through a full 180$^{\\circ}$ in position angle around the nucleus of Coma A, but there are no sharp boundaries in the surface brightness of the structures, corresponding to the edges of an ionization cone. It is possible for the emission line distributions to appear broader than the nominal  45-60$^{\\circ}$ cones predicted by the unified schemes if the cone axes are tilted towards the observer. However, in order to explain the emission line distributions in 3C171 and Coma A in this way, the cones would have to be tilted to such an extent that the observer's line of sight would lie within the cone and we would see the illuminating AGN directly. Clearly this is not the case, and it appears highly unlikely that the extended filaments away from the radio axis are photoionized by a central source of ionizing photons.  The most plausible alternative to quasar illumination is ionization by the shocks associated with the expanding radio jets and lobes. The emission lines could be produced as the warm clouds cool behind the shock fronts or, alternatively, as a consequence of photoionization of precursor clouds by the ionizing photons produced in the cooling, shocked gas. In either case we would expect a close morphological association between the radio and optical structures, just as we observe in 3C171 and Coma A. By adapting equation 4.4 of Dopita and Sutherland (1996), and assuming a shock speed through the warm clouds of 200 km s$^{-1}$, we estimate that the rate of flow of warm ISM through the shocks would have to be at least 1.9$\\times$10$^4$ M$_{\\odot}$ yr$^{-1}$ for 3C171, and 3.2$\\times$10$^3$ M$_{\\odot}$ yr$^{-1}$ for Coma A, in order for the emission line luminosities of the nebulae as a whole to be produced entirely by shock ionization. Energetically, the shock ionization mechanism appears to be feasible in the sense that the total emission line luminosities of the sources are $<$10\\% of the bulk powers of the radio jets (Clark 1996, Clark \\etal\\ 1998)\\footnote{In order to derive this result we have scaled the results of Clark (1996), who considered only the emission line components along the radio axis, to the total emission line fluxes for the nebulae as a whole, as derived from our H$\\alpha$ images.}.    However, it is not possible to rule out some contribution to the ionization of the extended structures by a central photoionizing source. As discussed in the introduction, some radio sources with relativistic jets may not have powerful quasar nuclei. If this is the case, the narrow beams of radiation emitted by the jets could contribute to the ionization of the structures along the radio axis, although the ionization of the more extended filamentary structures would continue to be  be dominated by interactions with the rdaio lobes.  One further possibility is that the structures are photoionized by young stars associated with the filaments. This is supported by the presence of faint continuum structures associated with the H$\\alpha$ filaments (see Figure 1(c)), and the spectroscopic detection of excess UV continuum emission to the north and south of the nucleus along the radio axis above the level expected for the nebular continuum emitted by the warm gas (Clark 1996). Without further information on the nature and spectrum of the extended continuum structures it is difficult to test this model at the present time.  An open question for both 3C171 and Coma A is the extent to which the structures reflect the true distributions of ionized gas in the haloes of the host galaxies, and the extent to which the structures are distorted by their interaction with the radio components. It is possible for shock fronts to sweep up material into shell-like structures. However, given that the clouds are likely to be destroyed by hydrodynamical interactions with the fast, hot wind behind the shock fronts within a few shock crossing times (\\eg\\ Klein, McKeee \\& Collella 1994), and given also the presence of diffuse H$\\alpha$ emission well away from the radio structures in both Coma A and 3C171, it seems more plausible that these represent pre-existing gas structures. Cloud destruction by shocks may also lead to a relative absence of warm gas in the lobes, further enhancing the shell-like appearance of the emission line structures. In the case of Coma A it is likely that we are seeing the results of interactions between between the radio-emitting components and the gaseous remnants of mergers/interactions in a group of galaxies.  Clearly, detailed measurements of the kinematics, line ratios, and continuum spectra of the filamentary structures are required in order to resolve the outstanding issues concerning the physical state, ionization and origins of the warm gas."
    },
    "0002/astro-ph0002342_arXiv.txt": {
        "abstract": "We report the serendipitous discovery of a population of low mass, pre-main sequence stars (PMS) in the direction of the Wolf-Rayet/O-star binary system $\\gamma^{2}$ Vel and the Vela OB2 association.  We argue that $\\gamma^{2}$ Vel and the low mass stars are truly associated, are approximately coeval and that both are at distances between 360-490\\,pc, disagreeing at the 2$\\sigma$ level with the recent Hipparcos parallax of $\\gamma^{2}$ Vel, but consistent with older distance estimates. Our results clearly have implications for the physical parameters of the $\\gamma^{2}$ Vel system, but also offer an exciting opportunity to investigate the influence of high mass stars on the mass function and circumstellar disc lifetimes of their lower mass PMS siblings. ",
        "introduction": "As little as 15 years ago it was believed that high mass stars formed mainly in OB associations and that low mass stars formed mainly in T associations. In the last few years it has become increasingly recognized that the majority of low mass stars in the solar neighbourhood are in fact likely to have formed in OB associations. The mass function in OB associations may match popular log-normal parameterisations of the general field population such as those proposed by Miller \\& Scalo (1979), with many low mass stars formed for every O/B star. The main reason for this shift in paradigm is the discovery of numerous low mass pre-main sequence (PMS) stars in young OB associations by virtue of their high levels of X-ray activity (see for example Walter et al. 1994, 1999; Naylor \\& Fabian 1999, Preibisch \\& Zinnecker 1999 and references therein). Recently, some evidence has been found that these X-ray selected groups can be quite small, concentrated around just one or two high mass stars, which are themselves part of larger associations. Examples include $\\beta$ Cru (Feigelson \\& Lawson 1997), $\\sigma$ Ori (Walter, Wolk \\& Sherry 1998) and $\\eta$ Cha (Mamajek, Lawson \\& Feigelson 1999).   The investigation of these low mass PMS stars is of prime importance in establishing the form of the {\\em initial} mass function. Although OB associations are by definition unbound, they are young enough that dynamical effects such as mass segregation and preferential evaporation will not have occurred; and they are usually free of the heavy, variable extinction that plagues observations of active star forming regions. If the majority of low mass stars are born in OB associations it is crucial to establish the influence that high mass neighbours have on the formation and evolution of their lower mass siblings. The winds and ionizing radiation of hot stars could influence the mass function and circumstellar disc lifetimes of the lower mass stars, with implications for angular momentum evolution and planet formation.  These ideas have gained currency with the discovery of evaporating discs around PMS stars in the Orion nebula (McCaughrean \\& O'Dell 1996) and theoretical studies showing that discs could be ionized and evaporated by the UV radiation fields of O stars (Johnstone, Hollenbach \\& Bally 1998). In more extreme circumstances, nearby supernova explosions have been invoked both as a means of terminating low mass star formation (Walter et al. 1994) or triggering it (Preibisch \\& Zinnecker 1999).  In this paper we report the discovery, by X-ray selection, of a low mass stellar population that seems likely to be associated with the nearest example of a Wolf-Rayet (WR) star, $\\gamma^{2}$ Velorum (HD 68273, HIP 39953, WR11). Like about half of the $\\simeq 200$ galactic WR stars known, it is a binary system (WC8+O8) with an orbital period of 78.5 days and a massive, interacting stellar wind. There is currently some controversy concerning the distance to $\\gamma^{2}$ Vel, which impacts upon the deduced luminosities, masses and mass loss rates from the system. It is important to get these parameters right because, as the nearest WR, $\\gamma^{2}$ Vel is an extreme test of stellar evolution models and calibrates the absolute magnitudes of WR stars.  The Hipparcos parallax yields a distance of $258^{+41}_{-31}$\\,pc to $\\gamma^{2}$ Vel (Schaerer, Schmutz \\& Grenon 1997, van der Hucht et al. 1997), in marked contrast to previous distance estimates which place it at 350-450\\,pc. The larger distance is in better agreement with the mean distance to the Vela OB2 association ($410\\pm12$\\,pc), of which $\\gamma^{2}$ Vel is the most massive proper-motion member (de Zeeuw et al. 1999).  The presence of a low mass stellar association in the extreme environment around $\\gamma^{2}$ Velorum offers an excellent empirical test of the possible influence of winds and ionizing radiation on low mass stars and their discs. Additionally it gives us a chance to measure the distance to $\\gamma^{2}$ Vel by matching PMS isochrones to the low mass stars. ",
        "conclusions": "From our discussion there seem to be two possible scenarios. (1) That the PMS stars are approximately at the same distance and age as $\\gamma^{2}$ Vel, and that this distance places $\\gamma^{2}$ Vel within the Vela OB2 association at 360-490\\,pc.  (2) That the PMS stars are part of the Vela OB2 association, possibly surrounding $\\gamma^{1}$ Vel, but that $\\gamma^{2}$ Vel is an isolated foreground object with no surrounding low mass stars at a similar age. We believe that (1) is {\\em far} more plausible than (2) because of the dispersion in the Vela OB2 Hipparcos parallaxes and the likely association of $\\gamma^{1}$ and $\\gamma^{2}$ Vel. Recently, the idea that $\\gamma^{2}$ Vel could form in isolation without accompanying low mass stars has also been challenged by the near IR detection of a K-type PMS companion only 4.7 arcsec distant (Tokovinin et al. 1999).  If the low mass PMS stars we have found are truly in the vicinity of $\\gamma^{2}$ Vel, they represent an exciting opportunity to explore the influence of adjacent high mass loss stars and ionizing UV radiation fields on the mass function and circumstellar disc lifetimes of low mass stars. It will be interesting to compare the frequencies of T-Tauri discs around these stars with the frequencies found in T associations and OB associations with similar ages. The PMS stars in Fig.2 have masses, found from the D'Antona \\& Mazzitelli (1997) models, down to (an age dependent) mass of about 0.15\\msun. The mass function will be addressed when we have a better census of the association membership."
    },
    "0002/astro-ph0002518_arXiv.txt": {
        "abstract": "The results of an X-ray observation of the nearby core-dominated, super-luminal quasar 4C\\,+73.18 (1928+738) with ASCA are reported. Variations of the  1--10\\,keV flux on time-scales of the order of 10 hours were found while the 0.5--1\\,keV flux remained almost unchanged. The measured flux at 1\\,keV had increased by $\\sim 50\\%$ compared with previous observations 18 years ago. The quasar shows a curved spectrum in the 0.8--10\\,keV band, which can be modeled as a power law plus a soft component. Both, the spectrum and the uncorrelated flux variations indicate the presence of a soft X-ray excess over the extrapolation of the hard band power law, which might be the high-energy extension of the intense UV bump emission found in this quasar. The hard X-ray emission is dominated by the flat power law (photon index 1.1--1.5). The results are consistent with models in which the hard X-rays are emitted from the relativistic jet. An iron line emission as reported in a previous Ginga observation (Lawson \\& Turner 1997) was not detected. We suggest that the line detected by Ginga might be not associated with the quasar but, most likely, with a distant cluster of galaxies. There are indications for a weak emission line-like feature around 1.25\\,keV (1.63\\,keV in the quasar frame) in both, the SIS and GIS spectra. The energy is surprisingly close to that of an unexplained line found in the flat-spectrum quasar PKS\\,0637-752 by Yaqoob \\eta (1998). No evidence for the proposed binary black holes (Roos \\eta 1993) is found in the current X-ray data. ",
        "introduction": "4C\\,+73.18 (1928+738) is an extremely core-dominated, super-luminal quasar at a redshift of 0.302. It is a flat-spectrum radio source in the S\\,5 survey (K\\\"uhr \\eta 1981), which has been well studied with VLBI and shows unusual jet properties. On arc-second (kpc) scales, the source exhibits two-sided curved jets, lobes, and a dominant core (Johnston \\eta 1987, Hummel \\eta 1992, Murphy \\eta 1993). On pc scales, VLBI observations of the core reveal an one-sided jet, which exhibits apparent super-luminal motion with $v_{\\rm app}/c \\sim 4-7$ (Eckart \\eta 1985,  Witzel et al.\\ 1988). The jet is estimated to be aligned within $7^{\\circ}-12^{\\circ}$ to the line-of-sight (e.g.\\ Ghisellini \\eta 1993, Jiang \\eta 1998). No $\\gamma$-ray emission was detected by EGRET (Fichtel \\eta 1994). To explain the mis-alignment between the VLBI jet and the kpc-scale jet, Hummel \\eta (1992) suggested that the observed VLBI jet might be only one of the Doppler-boosted filaments. Interestingly, the most recent results from the space VLBI (VSOP) monitoring of 4C\\,+73.18 reveal substantial jet bending at a distance of 6.3\\,pc from the core, and dramatic temporal variations of the bend angle over a few months (Murphy et al.\\ 1999).  It is intriguing that the sub-milliarcsec jet exhibits ballistic motion along a sinusoidal curve with a period of about 3 years (Hummel \\eta\\ 1992). The wiggles over such short period were modeled by Roos \\eta (1993) as the orbital motion of a binary black hole system, which has a mass of the order of $10^8\\,M_{\\odot}$, a mass ratio $>0.1$, and a separation of $\\sim 10^{16}$\\,cm. No other observational evidence for the binary black hole model has been reported so far. X-ray observations, having the advantage of looking into the region closed to the black hole, provide a potential tool to find signatures of binary black holes.  In the X-ray band, 4C\\,+73.18 has been observed previously by {\\em Einstein} (Biermann et al.\\ 1981), {\\em EXOSAT} (Biermann \\eta 1992, Ghosh \\& Soundararajaperumal 1992), Ginga (Lawson \\& Turner 1997), and by ROSAT in the Survey (Voges et al.\\ 1999) and pointed observations with PSPC (Brunner \\eta 1994, and Sambruna 1997) and HRI, respectively. The measured spectra were generally described by a power law. Of particular interest is the detection of a significant iron K$_\\alpha$ emission line in the Ginga observation, making this object one of the few flat-spectrum radio quasars showing iron K$_\\alpha$ line. Furthermore, the ROSAT data seem to suggest an absorption column in excess of the Galactic value.  In this paper, we report on results of an X-ray observation of 4C\\,+73.18 in the 0.8--10\\,keV band performed by the ASCA satellite. The spectral and temporal analysis of the data are presented in \\S\\,2 and 3, respectively. A re-analysis of the archival ROSAT PSPC spectrum is also given (\\S\\,\\ref{spec:rosat}). We discuss the implications of the results and summarize the main conclusions in \\S\\,4 and \\S\\,5, respectively. Errors are quoted at 68\\% level throughout the paper, unless mentioned otherwise. ",
        "conclusions": "The flat-spectrum, super-luminal quasar 4C\\,+73.18 resembles the well studied, similar object 3C\\,273 in some basic properties in the soft-to-medium X-ray band. Rapid, small amplitude ($\\sim 25\\%$) flux variations were detected in the 1--10\\,keV band on a tentative time scale of the order of 10 hours, whereas no significant variations were found for the soft X-rays below 1\\,keV. The 0.8--10\\,keV spectrum exhibits a curvature, which can be modeled by a power law plus a soft component. Both, the X-ray spectrum and the uncorrelated flux variations provide strong evidence for the presence of a soft X-ray excess over the extrapolation of the hard band power law, which might be the high-energy extension of the UV bump emission. The underlying continuum is flat, with a power law photon index of 1.1--1.5, depending on the spectral model adopted for the soft component. The dominance of the flat power law over possible Seyfert-like emission, together with the short time scale variability, are naturally explained by beaming models for the jet, which is estimated to be observed at an inclination angle about 10 degree. No iron K$_{\\alpha}$ emission line is detected, in contrast to the significant line reported in a previous Ginga observation, which is most likely contamination from another source. A weak emission line-like feature around 1.25\\,keV (1.6\\,keV in the quasar frame) seems to be present in both the SIS and GIS spectra, though the significance is not high. The line energy is identical to that of a mysterious emission line in the flat-spectrum quasar PKS\\,0637-752 reported by Yaqoob \\eta (1998)."
    },
    "0002/astro-ph0002032_arXiv.txt": {
        "abstract": "We report 1405 MHz measurements of the flux density of the $\\approx 320$ year old supernova remnant Cassiopeia A, relative to the flux density of Cygnus A, made between 1995 and 1999.  When compared to measurements made between 1957 and 1976, we find that the rate at which Cassiopeia A has been fading at this and nearby frequencies has changed from $\\approx 0.9$ \\% yr$^{-1}$ in the 1960s to $\\approx 0.6 - 0.7$ \\% yr$^{-1}$ now.  Furthermore, we have collected from the literature measurements of this fading rate at lower (38 -- 300 MHz) and higher (7.8 -- 16.5 GHz) frequencies.  We show that the fading rate has dropped by a factor of $\\approx 3$ over the past 50 years at the lower frequencies, while remaining relatively constant at the higher frequencies, which is in agreement with the findings of others.  Our findings at 1405 MHz, in conjunction with a measurement of the fading rate at the nearby frequency of 927 MHz by Vinyajkin (1997), show an intermediate behavior at intermediate frequencies.  We also find that Cassiopeia A, as of $\\approx$ 1990, was fading at about the same rate, $\\approx 0.6 - 0.7$ \\% yr$^{-1}$, at all of these frequencies.  Future measurements are required to determine whether the fading rate will continue to decrease at the lower frequencies, or whether Cassiopeia A will now fade at a relatively constant rate at all of these frequencies. ",
        "introduction": "Cassiopeia A is thought to have exploded around $1681 \\pm 15$, based on radial velocity and proper motion measurements of 15 of the supernova remnant's high-velocity [N II]-emitting knots, also called fast-moving flocculi or FMFs (Fesen, Becker, \\& Goodrich 1988).  Indeed, Flamsteed (1725) observed a 6th magnitude star within 13$\\arcmin$ of Cassiopeia A's location in 1680 (Ashworth 1980) that later became a subject of debate when the star could not be found in the sky (Herschel 1798; Baily 1835).  Cassiopeia A was re-discovered in 1943 at 160 MHz by Reber (1944), although a re-analysis of a day-long strip chart recording published by Jansky (1935) suggests that he unknowingly detected it as early as 1932 at 20.5 MHz (Sullivan 1978).  Accurate measurements of Cassiopeia A's flux density - both absolute measurements and measurements relative to the flux density of the extragalactic radio source Cygnus A - began in 1949 at 81.5 MHz (Ryle \\& Elsmore 1951).  \\subsection{The Fading Rate of Cassiopeia A around 1965}  H\\\"ogbom \\& Shakeshaft (1961), with Ryle \\& Elsmore's (1951) 1949 measurement and two measurements of their own made in 1956 and 1960 also at 81.5 MHz, first showed that Cassiopeia A was fading; they measured a fading rate of $1.06 \\pm 0.14$ \\% yr$^{-1}$, which was later refined to $1.29 \\pm 0.08$ \\% yr$^{-1}$ with the addition of four measurements made between 1966 and 1969 also at 81.5 MHz (Scott, Shakeshaft, \\& Smith 1969). Using data spanning 1957 -- 1972, Baars \\& Hartsuijker (1972) measured significantly lower fading rates at 1420 and 3000 MHz, $0.89 \\pm 0.12$ and $0.92 \\pm 0.15$ \\% yr$^{-1}$, respectively; they first suggested that Cassiopeia A was fading at different rates at different frequencies.  By 1977, accurate fading rates had been determined at six different frequencies, spanning 81.5 to 9400 MHz.  Baars et al. (1977; see also Dent, Aller, \\& Olsen 1974) collected these measurements, which we re-list in Table 1 and plot in Figure 1, and determined the following empirical equation that describes the frequency dependence of Cassiopeia A's fading rate for the $\\approx$ 1965 epoch: \\begin{equation} -\\frac{100}{F_{\\nu}}\\frac{dF_{\\nu}}{dt} = 0.97(\\pm0.04)-0.30(\\pm0.04)\\log{\\nu_{GHz}}. \\end{equation} Using this equation and absolute measurements of Cassiopeia A's flux density made around 1965 over a frequency range that spans 22 MHz to 31 GHz, Baars et al. (1977) also determined an empirical equation that describes the absolute spectrum of Cassiopeia A for the 1965 epoch to an accuracy of $\\approx 2$\\%.  Empirical equations that describe the absolute spectra of Cygnus A and Taurus A, and a semi-absolute spectrum for Virgo A, were also determined.  \\subsection{The Decreasing Fading Rate of Cassiopeia A at Low Radio Frequencies (38 -- 300 MHz)}  Equation (1) was first challenged by Erickson \\& Perley (1975) who's 1974 measurement of the flux density of Cassiopeia A at 38 MHz was brighter than the prediction of Baars et al. (1977) at the 3.5 $\\sigma$ level.  Subsequent measurements at 38 MHz by Read (1977a,b), Walczowski \\& Smith (1985), Rees (1990), and Vinyajkin (1997) confirmed that the fading rate had decreased from $1.9 \\pm 0.5$ \\% yr$^{-1}$ (determined using only the 1955 -- 1966 measurements; Read 1977a) to $0.66 \\pm 0.17$ \\% yr$^{-1}$ (determined using all of the measurements through 1995; Vinyajkin 1997).\\footnote{Rees (1990) suggested that the measurements at 38 MHz, as of 1987, could be treated as being consistent with a constant fading rate of $\\sim 0.8$ \\% yr$^{-1}$, particularly if one ignored the 1949 -- 1969 measurements of Ryle \\& Elsmore (1951), H\\\"ogbom \\& Shakeshaft (1961), and Scott, Shakeshaft, \\& Smith (1969), which imply a significantly higher fading rate, $1.29 \\pm 0.08$ \\% yr$^{-1}$, at the nearby frequency of 81.5 MHz.  However, Rees (1990) did not rule out the possibility of a changing fading rate at 38 MHz, and indeed, the later measurement of this fading rate by Vinyajkin (1997) suggests that it has changed at the $\\approx 2.3$ $\\sigma$ confidence level.}  Similar and more accurate results have been found at 81.5 MHz.  As stated above Scott, Shakeshaft, \\& Smith (1969) found a 1949 -- 1969 fading rate of $1.29 \\pm 0.08$ \\% yr$^{-1}$ at this frequency.  Hook, Duffett-Smith, \\& Shakeshaft (1992) found a significantly lower fading rate of $0.92 \\pm 0.16$ \\% yr$^{-1}$ when they included their 1989 measurement, and a fading rate of $\\approx 0.63$ \\% yr$^{-1}$ when they considered only the measurements made after 1965.  Agafonov (1996) found similar, although slightly higher values:  $1.25 \\pm 0.06$ \\% yr$^{-1}$ (1949 -- 1985) and $\\sim 0.8$ \\% yr$^{-1}$ (1973 -- 1985).  Consistent measurements have also been made at 102 MHz (Agafonov 1996), 151 -- 152 MHz (Read 1977a; Agafonov 1996; Vinyajkin 1997), and 290 -- 300 MHz (Baars \\& Hartsuijker 1972; Vinyajkin 1997).  We plot the low frequency (38 -- 300 MHz) measurements of Cassiopeia A's fading rate, as well as the intervals of time over which these measurements were made, in Figure 2, and we list this information in Table 2.  Clearly, as is suggested by Hook, Duffett-Smith, \\& Shakeshaft (1992), and advocated by Agafonov (1996), the rate at which Cassiopeia A is fading at low frequencies has decreased over the past 50 years.\\footnote{Hook, Duffett-Smith, \\& Shakeshaft (1992) suggested that instrumental uncertainties and/or flaring of Cassiopeia A's radio emission might systematically affect measurements of Cassiopeia A's fading rate at these low frequencies.  However, they made arguments to the contrary, and Agafonov (1996) provided evidence to the contrary.  Consequently, we do not further pursue these scenarios in this paper; instead, we refer the interested reader to these papers.}  We find the fading rate to be decreasing by $\\approx 2$ \\% yr$^{-1}$ per century at these frequencies.  \\subsection{The Constant Fading Rate of Cassiopeia A at High Radio Frequencies (7.8 -- 16.5 GHz)}  However, different conclusions have been reached at significantly higher frequencies.  O'Sullivan \\& Green (1999) compare four measurements of Cassiopeia A's relative brightness to Cygnus A that they made in 1994 and 1995 at 13.5, 15.5, and 16.5 GHz to the predictions of Baars et al. (1977).  The measurements of O'Sullivan \\& Green (1999) are perfectly consistent with the $\\approx 0.6$ \\% yr$^{-1}$ predictions of Equation (1),\\footnote{Technically, Equation (1) is an extrapolation at frequencies greater than 9.4 GHz; however, this is not a large extrapolation.} suggesting that the rate at which Cassiopeia A fades has not changed significantly over the course of the last half century at these high frequencies.  We plot the high frequency (7.8 -- 16.5 GHz) measurements of Cassiopeia A's fading rate in Figure 3, and we list this information in Table 3.  These apparently contradictory behaviors at low and high radio frequencies suggests that the rate at which Cassiopeia A fades is changing in a frequency-dependent way:  the fading rate is decreasing at the low frequencies (38 -- 300 MHz) and is relatively constant at the high frequencies (7.8 -- 16.5 GHz).  In this paper, we investigate and confirm this trend at the intermediate frequency of 1405 MHz with measurements of the relative flux density of Cassiopeia A to Cygnus A that we made between 1995 and 1999 with the 40-foot telescope at the National Radio Astronomy Observatory in Green Bank, WV.  We analyze these measurements in \\S 2; we draw conclusions in \\S 3. ",
        "conclusions": "We plot the intermediate frequency (927 -- 3060 MHz) measurements of Cassiopeia A's fading rate in Figure 5, and we list this information in Table 5.  When compared to measurements of the fading rate between 1957 and 1976, our measurement at 1405 MHz, in conjunction with a recent measurement at the nearby frequency of 927 MHz by Vinyajkin (1997), show that the fading rate is decreasing at a rate that is intermediate to the rates measured at the lower (38 -- 300 MHz) and the higher (7.8 -- 16.5 GHz) frequencies.  We find the fading rate to be decreasing by $\\approx 1$ \\% yr$^{-1}$ per century at these intermediate frequencies.  Furthermore, we find that around 1990, Cassiopeia A was fading at about the same rate, $\\approx 0.6 - 0.7$ \\% yr$^{-1}$, at all of these frequencies.  The next decade of observations should reveal whether the fading rate will continue to decrease at the lower frequencies, or whether Cassiopeia A will now fade at a relatively constant rate at all of these frequencies."
    },
    "0002/astro-ph0002204_arXiv.txt": {
        "abstract": "{For a universe presently dominated by static or dynamic vacuum energy, cosmological constant (LCDM) or quintessence (QCDM), we calculate the asymptotic collapsed mass fraction as function of the present ratio of vacuum energy to clustered mass, $\\omq/\\omm$. Identifying these collapsed fractions as anthropic probabilities, we find the present ratio $\\omq/\\omm \\sim 2$ to be reasonably likely in LCDM, and very likely in QCDM.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002210_arXiv.txt": {
        "abstract": "Observations are presented showing the doublet \\civww\\ absorption lines superimposed on the \\civ\\ emission in the radio galaxy 0943--242. Within the errors, the redshift of the absorption system that has a column density of $\\nciv = 10^{14.5 \\pm 0.1} \\,\\cms$ coincides with that of the deep \\lya\\ absorption trough observed by R\\\"ottgering et~al. (1995). The gas seen in absorption has a resolved spatial extent of at least 13\\,kpc (the size of the extended emission line region).  We first model the absorption and emission gas as co-spatial components with the same metallicity and degree of excitation. Using the information provided by the emission and absorption line ratios of \\civ\\ and \\lya, we find that the observed quantities are incompatible with photoionization or collisional ionization of cloudlets with uniform properties.  We therefore reject the possibility that the absorption and emission phases are co-spatial and favour the explanation that the absorption gas has low metallicity and is located further away from the host galaxy (than the emission line gas). The larger size considered for the outer halo makes plausible the proposed metallicity drop relative to the inner emission gas. In absence of confining pressure comparable to that of the emission gas, the outer halo of 0943--242 is considered to have a very low density allowing the metagalactic ionizing radiation to keep it higly ionized. In other radio galaxies where the jet has pressurized the outer halo, the same gas would be seen in emission (since the emissivity scales as $n_H^2$) and not in absorption as a result of the lower filling factor of the denser condensations. This would explain the anticorrelation found by Ojik et~al. (1997) between \\lya\\ emission sizes (or radio jet sizes) and the observation (or not) of \\hi\\ in absorption. The estimated low metallicity for the absorption gas in 0943--242 ($Z \\sim 0.01 \\zsol$) and its proposed location --outer halo outside the radio cocoon-- suggest that its existence preceeds the observed AGN phase and is a vestige of the initial starburst at the onset of formation of the parent galaxy. ",
        "introduction": "Very high redshift ($z>2$) radio galaxies (hereafter HZRG) show emission lines of varying degree of excitation. In virtually all objects, the \\lya\\ line is the strongest and is usually accompanied by high excitation lines of \\civw, \\ciiiw, \\heiiuw\\ and, at times, \\nvw\\ (R\\\"ottgering et~al. 1997 and references therein).  An important characteristic of the emision gas is its spatial scale.  The sizes of the \\lya\\ emission region range from $15$ to 120\\,kpc (van~Ojik \\etal\\ 1997).  Most ground work on HZRG is performed at rather low resolution ($\\sim 20$\\AA) to maximize the probability of line detection and the S/N. However a very potent discovery was made by van~Ojik \\etal\\ (1997, hereafter vO97) at much higher resolution, that of extended \\hi\\ {\\it absorption} gas. In effect, out of 18 HZRG spectra taken at the unusually high resolution of $\\simeq 1.5$--3\\AA, vO97 found --in 60\\% of the objects-- deep absorption troughs superimposed on the Ly$\\alpha$ emission profiles.  Furthermore, out of the 10 radio galaxies smaller than 50\\,kpc, strong \\hi\\ absorption is found in 9 of them. The absorption gas appears to have a covering factor near unity over very large scales, namely as large as the underlying emission gas.  The current paper addresses the problem of the ionization state of both the absorption and the emission gas as well as the interconnection between the two. The main justifications behind this work are the following: HZRG are probably the progenitors of the massive central cluster galaxies (Pentericci et~al. 1999) and as such are an important means by which we can study large ellipticals and their environment at such high redshift, a time not so long after, or even during their formation. Furthermore, the extended gas as detected in \\civ\\ (see below) is chemically enriched and therefore represents debris of past intense stellar formation periods and is interesting to study in their own right. What is the fate of such gas?  How quickly has the enrichment of this large scale gas proceeded? Will this gas be heated up into a hot wind and enrich the intergalactic X-ray gas in cluster of galaxies? Will it on the contrary condense into sheets or condensations? A better understanding of the various gaseous phases which co-exist in high redshift objects would help anwering these questions.  To determine the physical conditions of the absorption gas, new observations were carried out at the wavelength of \\civ\\ and \\heii\\ in 0943--242, the first radio galaxy reported to show large scale absorption troughs (R\\\"ottgering et~al. 1995, hereafter RO95).  The new spectrum shows the \\civ\\ absorption doublet at the same redshift\\footnote{We will distinguish between absorption and emission redshifts using subscripts, as in $z_a$ and $z_e$, respectively.} $z_{a}$ as the \\lya\\ absorption trough (RO95). Clearly and surprisingly the gas in absorption is highly ionized and probably of comparable excitation to the gas seen in emission.   This paper is structured as follows. We first present observations which show \\civ\\ in absorption in 0943--242 (Sect.\\,~\\ref{observ}). In Sect.\\,~\\ref{hypo} we derive a ratio (\\gamo) relating the observed emission and absorption quantities which depends somewhat on the ionization fraction of H but not explicitely on the C/H metallicity ratio.  At first, we postulate that the emission and absorption gas components are co-spatial and share the same excitation mechanism and physical conditions and proceed to model \\gamo\\ with a one-zone equilibrium photoionization model. We improve on the model using a stratified photoionized slab. As the observed ratio cannot be reproduced even in the case of collisional ionization, we discuss in Sect.~\\ref{intergam} two alternative interpretations of this significant discrepancy. We demonstrate the many advantages of the winning scenario in which the absorption gas is further out and of much lower density, pressure and metallicity than the emission gas. ",
        "conclusions": "\\label{conclusions}  The detection of \\civ\\ absorption in radio galaxy 0943--242 at the same redshift as the deep \\lya\\ trough observed by RO95 demonstrates that the detected absorption gas is highly ionized.  Having assumed that the \\hi\\ and \\civ\\ columns measured from the Voigt profile fitting were representative of the dominant gas phase (by mass) in the outer halo, we have effectively ruled out that the absorption and emission gas occupy the same position in 0943--242.  We subsequently reassessed the picture proposed by vO97 in which both the large scale emission gas and the absorption gas were of comparable density ($n_H \\sim 100\\,\\cmc$).  In the former picture, the absorption gas was believed to lie outside the AGN ionization bicone (see their Fig.~11 in vO97). To ionize the gas to such a degree without using the AGN flux is problematic.  We have proposed an alternative picture in which the absorption gas is of very low metallicity and lies far away (in the outer halo) from the inner pressurized radio jet cocoon. Since in this new scheme the density of the absorption gas is expected to be very low, the metagalactic background radiation now suffices to photoionize it.  Furthermore, the structure of the absorption gas is now drastically simplified since we do not need over $\\sim 10^{10}$ condensations of size $\\sim 1$pc and density $\\sim 10^2\\,\\cmc$ to reach a covering factor close to unity. We can now reach similarly high covering factor using a single or few shells of very low density which have a volume filling factor close to unity (assuming a density of $\\sim 10^{-2.5}\\,\\cmc$).  It appears to us that the low metallicity inferred ($Z \\simeq 0.01 \\zsol$) and the proposed location of the absorption gas in 0943--242 --outside the radio cocoon, in an outer halo which is  seen in emission in other radio galaxies (as in 1243+036)-- strongly suggest that the absorbers' existence precedes the observed AGN phase.  Unless this non-primordial gas has been enriched by still undetected pop III stars, we consider that it more likely corresponds to a vestige gas phase expelled from the parent galaxy during the initial starburst at the onset of its formation.  If the \\civ\\ doublet was detected in absorption in other radio galaxies with deep \\lya\\ absorption troughs, there are many aspects which would be worth studying.  For instance, how uniform is the excitation of the absorption gas across the region over which it is detected?  Is a single phase sufficient? This could be tested by an attempt to detect absorption troughs of \\mgiiw\\ or imaging the troughs in \\civ\\ with an integral field spectrograph on an 8-m class telescope.  How different is the metallicity of the absorption gas in the other radio galaxies?  The information gathered could then be used to infer the enrichment history of the outer halo gas which surrounds HZRG."
    },
    "0002/astro-ph0002026_arXiv.txt": {
        "abstract": "To explain the origin of Herbig Ae/Be stars activity, it has been recently proposed that strong mass-losses trigger rotational instabilities in the envelope of fast rotating stars. The kinetic energy transferred to turbulent motions would then be the energy source of the active phenomena observed in the outer atmosphere of Herbig Ae/Be stars (Vigneron et al. 1990; Ligni\\`eres et al. 1996).  In this paper, we present a one-dimensional model of angular momentum transport which allows to estimate the degree of differential rotation induced by mass-loss. Gradients of angular velocity are very close to $- 2 \\Omega / R$ ($\\Omega$ being the surface rotation rate and $R$ the stellar radius). For strong mass-loss, this process occurs in a short time scale as compared to other processes of angular momentum transport. Application of existing stability criteria indicates that rotational instabilities should develop for fast rotating star. Thus, in fast rotating stars with strong winds, shear instabilities are expected to develop and to generate subphotospheric turbulent motions. Albeit very simple, this model gives strong support to the assumption made by Vigneron et al. 1990 and Ligni\\`eres et al. 1996. ",
        "introduction": "Although far from complete, the overall picture describing the angular momentum evolution of solar type stars is well established. Despite modest mass-loss rates, $\\dot{M} = 10^{-14} {\\rm M}_{\\sun} {\\rm yr}^{-1}$ for the sun, magnetised stellar winds strongly brake the rotation of the star. The decrease of the angular velocity then reacts back on the magnetic field by reducing the efficiency of the dynamo process. This picture cannot be applied as such to early-type stars. First of all, the history of stellar rotation at intermediate and high masses is generally still poorly known; besides, the presence of magnetic fields in early-type stars has not been established, except for some categories of chemically peculiar stars, as well as for a couple of particular cases (Donati et al. 1997; Henrichs et al. 2000). On the other hand, we know that early type stars can experience strong angular momentum losses as stellar winds with very large mass loss rate have been observed ($\\dot{M} \\approx 10^{-5} - 10^{-8} {\\rm M}_{\\sun} {\\rm yr}^{-1}$). In this paper, we investigate how such angular momentum losses will affect the stellar rotation, assuming magnetic fields are not dynamically relevant.  This question has not yet received much attention although it might have important consequences for stellar structure and evolution as well as for the understanding of early-type stars activity. In the context of stellar evolution, the effect of mass-loss on rotation has to be investigated since the rotation strongly influences the stellar structure. However, current models of stellar evolution with rotation do not take this effect into account (Talon et al. 1997, Denissenkov et al. 1999). As explained below, we have been confronted to the same question while investigating the origin of the very strong activity of Herbig Ae/Be stars, a class of pre-main-sequence stars with masses ranging from $2$ to $5 {\\rm M}_{\\sun}$.  A significant fraction of these objects is known to possess extended chromospheres, winds, and to show high levels of spectral variability. In addition, the presence of magnetic fields, first suggested by the rotational modulation of certain spectral lines (Catala et al. 1986), has been recently supported by a direct detection at the surface of the Herbig Ae star, HD 104237 (Donati et al. 1997). Detailed estimates of the non-radiative heating in the outer atmosphere of Herbig Ae/Be stars (Catala 1989, Bouret et al. 1998) compared to observational constraints on the available energy sources strongly suggest that the rotation of the star is the only energy source capable of powering such activity (see discussion in F.Ligni\\`eres et al. 1996). This led Vigneron et al. (1990) to propose a scenario whereby the braking torque exerted by the stellar wind forces turbulent motions in a differentially rotating layer below the stellar surface. Then, by invoking an analogy with stellar convection zones, these turbulent motions could generate a magnetic field which would transfer and dissipate the turbulent kinetic energy into the outer layers of the star.  Unlike Herbig stars, there is at present no observational evidence of non-radiative energy input in the atmosphere of OB stars. Their radiatively-driven winds do not require any additional acceleration mechanism and an eventual non-radiative heating would be very difficult to detect because most lines are saturated. However, they show various forms of spectral variability. Recent observations seem to indicate that these phenomena are not due to an intrinsic variability of the wind but are instead caused by co-rotating features on the stellar surface (Massa et al. 1995). As proposed in the literature (Howarth et al. 1995; Kaper et al. 1996), these corotating features could be due to a magnetic structuration of the wind. Since OB stars rotate fast and possess strong winds, the Vigneron et al. scenario might also explain these phenomena. The recent direct detection of a magnetic field in an early Be star, $\\beta$~Cep, adds some credit to this hypothesis (Henrichs et al. 2000).  In this paper, we shall investigate the starting assumption of the Vigneron et al. scenario namely that the braking torque of an unmagnetised wind generates strong enough velocity shear to trigger an instability in the subphotospheric layers. This is a crucial step in the scenario since the onset of the instability allows to transfer kinetic energy from rotational motions to turbulent motions.  The possible connection between mass-loss and instability had been already suggested by Schatzman (1981). However, the onset of the instability is not really considered in this study since it is assumed from the start that the wind driven angular momentum losses induce a turbulent flux of angular momentum. Here, we propose a simple model of angular momentum transport by a purely radial mass flux in order to estimate the angular velocity gradients induced by mass loss. Then, the stability of these gradients is studied according to existing stability criteria. Note that, as we neglect latitudinal flows, our model is best regarded as an equatorial model.  In the absence of magnetic fields, the braking effect of a stellar wind is simply due to the fact that matter going away from the rotation axis has to slow down to preserve its angular momentum. Thus, for the star's surface to be significantly braked, fluid parcels coming from deep layers inside the star have to reach the surface. Since radial velocities induced by mass-loss are very small deep inside the star, this is expected to take a relatively long time, not very different from the mass-loss time scale. By contrast, we shall see that the formation of unstable angular velocity gradients near the surface take place in a very short time, much smaller than the braking time scale.  The paper is organised as follows: first, we estimate the time scale characterising the braking of the stellar surface and relate it to the mass loss time scale (Sect. 2). Then, we show that radial outflows in stellar envelopes generate differential rotation and we estimate the time scale of this process (Sect. 3). The stability of these angular velocity gradients is considered (Sect. 4) and the results are summarised and discussed (Sect. 5). ",
        "conclusions": "We have studied the effect of mass-loss on the rotation of stellar envelopes in the simplified context of an unmagnetised spherically symmetric outflow. In the same way as an inhomogeneous distribution of a passive scalar tends to be smoothed out in a expanding flows, the radial gradients of specific angular momentum are smoothed out by radial expansion in stellar envelopes. This process becomes more and more effective as one approaches the surface because expansion becomes stronger and stronger. As shown in Figs 4 and 5, a profile of uniform angular momentum, $\\Omega \\propto 1/r^2$ is rapidly set-up in the outermost layers of the star and then pervades towards the interior on much larger time scales. These time scales being proportional to the mass-loss time scale, they vary very much from stars to stars. For a typical Herbig star ($\\dot{M} = 10^{-8}  {\\rm M}_{\\sun} {\\rm yr}^{-1}$), an angular velocity gradient close to $-2 {\\Omega}/R$ appears in one month. By contrast, for a solar-type mass-loss rate, $\\dot{M} = 10^{-14} {\\rm M}_{\\sun} {\\rm yr}^{-1}$, it would be $10^6$ times longer to reach the same level of differential rotation. According to existing stability criteria, the uniform angular momentum profile is subjected to shear instabilities provided the rotation period is shorter than $18.4 \\; g_{\\sun}/g$ days.  The present one-dimensional model is admittedly not realistic at least because latitudinal variations occur in rotating stars and give rise to a meridional circulation. Nevertheless, as already mentioned, neglecting latitudinal flows may be justified in the equatorial plane. In the following we discuss to which extent the neglected processes can prevent the formation of turbulent differentially rotating layers below the surface of rapidly rotating mass-losing stars. We first consider the angular momentum transport and then the stability problem.  Latitudinal variations could be inherent to the mass-loss mechanism as proposed for radiatively driven wind emitted from fast rotating stars (Owocki \\& Gayley 1997). Such variations are likely to generate a meridional circulation (Maeder 1999) but, at the latitudes where outflows occur, we still expect radial expansion to play a major role in shaping the angular velocity profile.  The Eddington-Sweet circulation driven by departure from sphericity operates on time scale larger than the Kelvin-Helmholtz time scale of the star. This time scale is of the order of the order of $10^7$ years for a $2 M_{\\sun}$ Herbig stars. Consequently the Eddington-Sweet circulation is unlikely to prevent the formation of the differentially rotating subphotospheric layer for strong mass-loss rate.  Turbulent motions like those generated in thermal convection zones might prevent the formation of the differentially rotating layer. In radiative envelopes however, it is not clear whether turbulent motions (not generated by the mass-loss process) would be vigorous enough.  Strong magnetic fields could in principle prevent the formation of these gradients. Note that this objection is not relevant in the context of the Vigneron et al. scenario where the shear layer is first produced by an unmagnetised wind.  Finally, we also neglected the effect of pre-main-sequence contraction. This is justified because the negative radial velocities associated with the contraction are much smaller than the positive radial velocities induced by mass-loss in the vicinity of the stellar surface.  In what concerns the stability of the uniform angular momentum profile, one has to remind that the stability criterion results from an extrapolation of laboratory experiments results. Despite recent numerical simulations which contradict the validity of this extrapolation for a differential rotation stable with respect to the Rayleigh-Taylor instability (Balbus et al. 1996), analysis of the experimental results reveal that for large values of the Reynolds number the properties of the turbulent flow do not depend on its stability or instability with respect to the Rayleigh-Taylor criterion (Richard \\& Zahn 1999). The solution of this debate would have to wait numerical simulations with higher resolution or specifically designed laboratory experiments.  The fact that the shear layer is embedded in an expanding flow may also affect the stability because expansion is known to suppress the turbulence. For example, this phenomenon is observed in numerical simulations of turbulent convection near the solar surface (Stein \\& Nordlund 1998).  However, the growth time scale of the instability being of the order of the rotation period ($\\approx 1$ day for Herbig stars), we do not expect the expansion to be rapid enough to suppress shear turbulence.  We conclude that the present model of angular momentum advection by a radial flow supports the assumption made by Vigneron et al. (1990) and F. Ligni\\`eres et al. (1996) that unmagnetised winds tend to force an unstable differential rotation in the subphotospheric layers of Herbig Ae/Be stars. Note that this work is part of an ongoing effort to assess the viability of the Vigneron et al. scenario. F. Ligni\\`eres et al (1996) and F.Ligni\\`eres et al. (1998) have already studied the step further, once turbulent motions are generated at the top of the radiative envelope, their work pointing towards the formation and inwards expansion of a differentially rotating turbulent layer. In addition, three-dimensional numerical simulations are being performed to investigate whether magnetic fields can be generated in such a layer."
    },
    "0002/astro-ph0002356_arXiv.txt": {
        "abstract": " ",
        "introduction": "The main purpose of the Iron Project (IP; Hummer \\etal 1993) is the continuing development of relativistic methods for the calculations of atomic data for electron impact excitation and radiative transitions in iron and iron-peak elements. Its forerunner, the Opacity Project (OP; Seaton \\etal 1994; The Opacity Project Team 1995), was concerned with the calculation of radiative parameters for astrophysically abundant elements, oscillator strengths and photoionization cross sections, leading to a re-calculation of new stellar opacities (Seaton \\etal 1994). The OP work, based on the non-relativistic formulation of the close coupling approximation using the R-matrix method (Seaton 1987, Berrington \\etal 1987), was carried out in LS coupling, neglecting relativistic fine structure that is not crucial in the calculation of mean plasma opacities. Also, collisional processes were not considered under the OP. The IP collaboration seeks to address both of these factors, and with particular reference to iron and iron-peak elements. The collaboration involves members from six countries: Canada, France, Germany, UK, US, and Venezuela.  The relativistic extension of the R-matrix method is based on the Breit-Pauli approximation (Berrington \\etal 1995). Collisional and radiative processes may both be considered. However, the computational requirements for the Breit-Pauli R-matrix (hereafter BPRM) calculations can be orders of magnitude more intensive than non-relativistic calculations. Nonetheless, a large body of atomic data has been obtained and published in a continuing series under the title {\\bf Atomic Data from the Iron Project} in {\\it Astronomy and Astrophysics Supplement Series}, with 43 publications at present. A list of the IP publications and related information may be obtained from the author's Website: www.astronomy.ohio-state.edu/~pradhan.  The earlier phases of the Iron Project dealt with (A) fine structure transitions among low-lying levels of the ground configuration of interest in Infrared (IR) astronomy, particularly the observations from the Infrared Space Observatory, and (B) excitation of the large number of levels in multiply ionized iron ions (with n = 2,3 open shell electrons, i.e. Fe~VII -- Fe~XXIV) of interest in the UV and EUV, particularly for the Solar and Heliospheric Observatory (SOHO), the Extreme Ultraviolet Explorer (EUVE), and Far Ultraviolet Spectroscopic Explorer (FUSE). In addition, the IP data for the low ionization stages of iron (Fe~I -- Fe~VI) is of particular interest in the analysis of optical and IR observations from ground based observatories. In the present review, we describe the IP work within the context of applications to X-ray spectroscopy, where ongoing calculations on collisional and radiative data for H-like Fe~XXVI, He-like Fe~XXV, and Ne-like Fe~XVII are of special interest.  The sections of this review are organised as follows: 1. Theoretical, 2. collisional, 3. radiative, 4. collisional-radiative modeling of X-ray spectra, 5. atomic data, and 6. Discussion and conclusion. ",
        "conclusions": "An overview of the work under the Iron Project collaboration was presented. Its special relevance to X-ray astronomy was pointed out since the IP, and related work, primarily aims to study the dominant atomic processes in plasmas, and to compute extensive and accurate set of atomic data for electron impact excitation, photoionization, recombination, and transition probabilities of iron and iron-peak elements. The importance of coupled-channel calculations was emphasized, in particular the role of autoionizing resonances in atomic phenomena. (A new project RMaX, a part of IP focused on X-ray spectroscopy, is described by K.A. Berrington in this review).  During the discussion, a question was raised regarding the resonances in Fe~XVII collision strengths (e.g. Fig. 1), and it was mentioned that new experimental measurments appear not to show the expected rapid variations in cross sections. A possible explanation may be that there are numerous narrow resonances in the entire near-threshold region, without a clearly discernible background or energy gap. The measured cross sections are averages over the resonances corresponding to the experimental beam-width. These averaged cross sections themselves may not exhibit sharp variations, unlike more highly charged He-like ions where the the non-resonant background and the resonance complexes are well separated in energy (e.g. He-like Ti~XXI, Zhang and Pradhan 1993).   \\def\\amp{{\\it Adv. At. Molec. Phys.}\\ } \\def\\apj{{\\it Astrophys. J.}\\ } \\def\\apjs{{\\it Astrophys. J. Suppl. Ser.}\\ } \\def\\apjl{{\\it Astrophys. J. (Letters)}\\ } \\def\\aj{{\\it Astron. J.}\\ } \\def\\aa{{\\it Astron. Astrophys.}\\ } \\def\\aasup{{\\it Astron. Astrophys. Suppl.}\\ } \\def\\adndt{{\\it At. Data Nucl. Data Tables}\\ } \\def\\cpc{{\\it Comput. Phys. Commun.}\\ } \\def\\jqsrt{{\\it J. Quant. Spectrosc. Radiat. Transfer}\\ } \\def\\jpb{{\\it Journal Of Physics B}\\ } \\def\\pasp{{\\it Pub. Astron. Soc. Pacific}\\ } \\def\\mn{{\\it Mon. Not. R. astr. Soc.}\\ } \\def\\pra{{\\it Physical Review A}\\ } \\def\\prl{{\\it Physical Review Letters}\\ } \\def\\zpds{{\\it Z. Phys. D Suppl.}\\ } \\def\\adndt{Atomic Data And Nuclear Data Tables}"
    },
    "0002/astro-ph0002483_arXiv.txt": {
        "abstract": "Models are presented for the polytropic equation of state of self-gravitating, quiescent interstellar gas clouds. A detailed analysis, including chemistry, thermal balance, and radiative transfer, is performed for the physical state of the gas as a function of density, metallicity, velocity field, and background radiation field. It is found that the stiffness of the equation of state strongly depends on all these physical parameters, and the adiabatic index varies between $\\sim 0.2-1.4$. The implications for star formation, in particular at high redshift and in starburst galaxies, and the initial stellar mass function are discussed. ",
        "introduction": "A fundamental problem in the theory of star formation is the physical structure of molecular gas clouds, such as  dense cores and Bok globules, from which (low-mass) stars are formed. Early work on the thermal and chemical balance of interstellar clouds has been performed by de Jong, Dalgarno \\& Boland (1980) and Falgarone \\& Puget (1985). A framework for active regions such as Orion has been developed by Tielens \\& Hollenbach (1985). The above work has been extended and refined by many authors (see Hollenbach \\& Tielens 1999 for a review), driven to a large extent by recent advances in observational techniques.  Examples of such observational progress include the catalog of 248 optically selected Bok globules of Clemens \\& Barvainis (1988), the NH$_3$ survey of isolated cores in Taurus by Benson \\& Myers (1989, and references therein), the measurement of non-thermal line widths in regions surrounding ammonia cores (Fuller \\& Myers 1992, Goodman et al.\\ 1998), and the use of spectroscopy of various atoms and molecules as temperature and density diagnostics of different types of interstellar clouds as summarized in van Dishoeck (1997).  Despite this wealth of observational data, the nature of their equation of state (EOS) remains a major theoretical problem that concerns the stability and collapse of these molecular clouds  The stiffness of the EOS can be largely responsible for the resulting density probability function of interstellar gas in the turbulent ISM (Scalo et al.\\ (1998) and Passot \\& Vazquez-Semadeni 1998). In particular, the value of gamma affects the densities that can be attained behind shocks (Vazquez-Semadeni, Passot \\& Pouquet 1996), the fraction of gas driven to high densities by turbulent interactions (c.f.\\ Scalo et al. 1998), and the stability of clouds being overrun by shocks (e.g.\\ Tohline, Bodenheimer \\& Christodoulou 1987; Foster \\& Boss 1996). In the present paper we concentrate on the effects of the EOS for the fragmentation of individual collapsing clouds. Before embarking on an analysis of the EOS of interstellar gas clouds, say of the form $P=P(\\rho ,T_{\\rm g},T_{\\rm d},V,B,I,C)$ for pressure $P$, density $\\rho$, gas temperatures $T_{\\rm g}$, dust temperature $T_{\\rm d}$, velocity $V$, magnetic field $B$, radiation intensity $I$, and chemical composition $C$, one should assess which of the possible physical quantities entering the EOS dominates in  the parameter ranges under study.  The importance of $\\rho$ and $T_{\\rm g}$ is self-explanatory. The presence of warm dust can play a role since it can heat the gas through gas-grain coupling and radiatively through optically thick molecular lines (Takahashi, Hollenbach \\& Silk 1983). The velocity field, even when it does not contribute to the turbulent pressure, strongly influences the optical depth of atomic and molecular cooling lines, and hence the thermal balance of the medium. The magnetic field plays an important role in the support of interstellar gas (McKee 1999). The presence of ultraviolet or other hard radiation sources strongly influences the chemical, thermal and ionization balance of the ambient medium, and hence the overall EOS. Finally, the chemical composition of the gas also depends on the metallicity since it is the abundance of atoms and molecules which influences the local cooling rate.  In the present work, a polytropic equation of state, $P=K\\rho^\\gamma$, is assumed to describe the physical state of the gas. The exponent $\\gamma$ is considered to be a function $\\gamma (T_{\\rm g},T_{\\rm d},V,I,C)$. That is, the magnetic pressure is assumed to be not dominant in the support of the model cloud. The radiation field $I$, in addition to the radiation emitted by the dust grains, is assumed to be given by the cosmic ray ionization rate. In other words, the visual extinction of the model cloud is assumed to be sufficient to shield the gas from any ultraviolet (stellar) sources.  We note that $\\gamma$ is the logarithmic derivative of $P$ with respect to $\\rho$. That is, since $P\\sim\\rho T$ it follows that $\\gamma = 1+{{d {\\rm log}T}\\over{d{\\rm log}\\rho}}$. Through implicit differentiation the above derivative is related to the logarithmic derivatives of the heating and cooling functions, which can be quite complex. It is this dependence of $\\gamma$ on the details of heating and cooling that is the main topic of the present paper.  In the remainder of the paper, a useful guide is the density dependence of the heating and cooling rate. If the density dependence of the heating rate is steeper (shallower) than that of the cooling rate, then $\\gamma$ is larger (smaller) than unity. A similar result holds for the kinetic temperature. Here, the cosmic ray ionization rate is always linear in the density and independent of temperature, whereas gas-grain heating and cooling is quadratic in the density and proportional to $T_{\\rm g}^{1/2}(T_{\\rm g}-T_{\\rm d})$. Optically thin, subthermally excited cooling lines have a quadratic density dependence which becomes linear in the thermalized (high density) limit and even effectively sub-linear when radiative trapping (optical depth effects) becomes important. Finally, collisional de-excitation rates are characterized by powers of the temperature on the order of $\\sim 0.5-1.5$, whereas excitation rates are modified by an exponential factor which is of the order of unity for kinetic temperatures larger than the pertaining excitation temperatures.  In general, our results are applicable mostly to dense and/or well shielded molecular clouds/cores where the pressure appears to be close to thermal, with only a small additional contribution to the line-widths from turbulent motions, and the ionization fraction is small (Goodman et al.\\ 1998, and references therein). That is, we assume that the clouds are outside the regime where the so-called Larson (1981) laws apply, which indicate an increase in the non-thermal line width with length scale.  One should note that polytropic models can also be extended into the regime where there is a significant contribution to the pressure from turbulence and magnetic fields. The polytropic temperature $T$ should then be interpreted as representing the total velocity dispersion $\\sigma$, i.e., $T\\propto\\sigma^2$ (Maloney 1988). A complete description of polytropic models, including composite ones for non-thermal pressure effects and bulk properties like mass, radius and density contrast, is presented in Curry \\& McKee (1999). The aim here is to study the thermal regime and to investigate the specific effects of metallicity, radiative transfer and internal sources on the stiffness of the resulting local EOS. ",
        "conclusions": "The polytropic equation of state of interstellar gas clouds has been computed for a range of physical parameters, and its relation to the IMF, high redshift star formation, and starburst galaxies has been investigated. A much needed extension of the results presented here lies in MHD, and particularly turbulence. The magnetic and velocity fields are important for pressure support away from the most centrally condensed part of a molecular cloud core, and radiative transfer effects in a turbulent medium are non-trivial. More to the point, the polytropic EOS indices computed here, with their relevance to the IMF as discussed in Section 4, are to be applied to a medium with a given cloud density probability distribution, as investigated in Scalo et al.\\ (1998). Such a cloud density probability distribution can only be derived from detailed hydrodynamic simulations which include magnetic pressure and Coriolis forces. Furthermore, $\\gamma$ can also affect, in a way still to be investigated, star formation and the underlying IMF in other scenarios like turbulence-induced condensation or shocked-cloud induced star formation. We would like to mention as well the work by Ogino, Tomisaka \\& Nakamura (1999) that shows that the mass accretion rate in their polytropic simulations of collapsing clouds depends a lot on whether gamma is taken as 0.8 or 1.2.  Extensions in terms of multiple polytropic models for the core and envelope of a molecular cloud, such as discussed in Curry \\& McKee (1999), are another avenue to explore for the bulk properties of interstellar clouds. Caveats in the approach presented here are grain mantle evaporation which can influence the abundance of H$_2$O and hence the heating rate in the presence of an IR background, uncertainties in the chemical reaction rates of neutral-neutral reactions at low temperatures, and details of the freeze-out onto grains of molecular species. We hope that through a synthesis of the approach adopted here and those of others mentioned above, one can address the question of cloud stability and collapse in a more general (cosmological) framework."
    },
    "0002/astro-ph0002095_arXiv.txt": {
        "abstract": "The {\\it Centre de Donn\\'ees astronomiques de Strasbourg} (CDS) provides homogeneous access to heterogeneous information of various origins: information about astronomical objects in {\\sc Simbad}; catalogs and observation logs in {\\sc VizieR} and in the catalogue service; reference images and overlays in {\\sc Aladin}; nomenclature in the {\\it Dictionary of Nomenclature}; Yellow Page services; the AstroGLU resource discovery tool; mirror copies of other reference services; and documentation. With the implementation of links between the CDS services, and with other on--line reference information, CDS has become a major hub in the rapidly evolving world of information retrieval in astronomy, developing efficient tools to help astronomers to navigate in the world-wide `Virtual Observatory' under construction, from data in the observatory archives to results published in journals.  The WWW interface to the CDS services is available at:  http://cdsweb.u-strasbg.fr/ ",
        "introduction": "The {\\it Centre de Donn\\'ees astronomiques de Strasbourg} (CDS) was founded in 1972 as the {\\it Centre de Donn\\'ees Stellaires}, and installed in Strasbourg as the result of an agreement between the INAG ({\\it Institut National d'Astronomie et de G\\'eophysique}), now INSU ({\\it Institut National des Sciences de l'Univers}), and {\\it Universit\\'e Louis Pasteur}. This was the outcome from the very prospective vision of Jean Delhaye, then Director of INAG, who anticipated the importance of computer-readable data. It was decided to found CDS as a part of the French research system, as a Data Center serving the international astronomical scientific community.  The objectives of CDS at its creation could be summarized as follows :  \\begin{itemize} \\item collect `useful' data about astronomical sources, in electronic form; \\item improve them by critical evaluation and combination; \\item distribute the results to the international community; \\item conduct research programs using the data collected. \\end{itemize}  Insertion of CDS into a research institution, the {\\it Observatoire astronomique de Strasbourg}, helps to maintain direct contacts with the evolution of astronomy, and with researchers' actual needs.  At the beginning, CDS was dealing with stellar data, aiming at the study of the galactic structure. In 1983, it was decided that {\\sc Simbad}, one of the two important CDS services at that time, would also deal with other galactic and extragalactic objets -- i.e., with all astronomical objects outside the Solar System. The CDS's name was changed to {\\it Centre de Donn\\'ees astronomiques de Strasbourg}, thus preserving the acronym which was already well known.  In recent years, research activities in astronomy have evolved significantly, with the very rapid development of on-line information at all levels, from observatory archives to results published in journals. The challenge is now to deal with {\\it information} more than with {\\it data}, which includes data, but also know-how about data, technical information about instruments, published results, compilations, etc.  The CDS goals can now be summarized as {\\bf collect, homogenize, distribute, and preserve astronomical information for the scientific use of the whole astronomical community}. This `mission statement' still contains all the drivers from the early CDS charter: dealing with electronic data, taking up an international role, developing expertise on astronomical data, having research as a goal.  All on-line CDS services can be accessed from the CDS home page on the World-Wide Web\\footnote{http://cdsweb.u-strasbg.fr/}. ",
        "conclusions": "The usage of the CDS services has undergone a revolution in the last few years, with the outcome of the World Wide Web, allowing easy access to on-line information, the integration of data and documentation, and navigation between distributed information. This means an explosion in the usage of the services, new functionalities, new concepts in the partnership of data centers and journals, since published information can now be considered as data (e.g., published tables are now usable like reference catalogues), and an increase in international partnership, to build up links and to define exchange standards. In parallel, the construction of the database contents remains a long-term activity, with increasing volume of information to deal with, and high standards of scientific and technical expertise needed in the value-added data center activities.  Navigation and links will certainly remain important keywords for the future, and the topic of interoperability is clearly emerging. One aspect is the construction of links between distributed services. Another one is the building of comprehensive information retrieval tools, as stressed in the AstroBrowse NASA initiative. To go further, one needs to be able to integrate the result of queries to heterogeneous services (ISAIA project, Hanisch \\cite{isaia}). In this context, the elaboration of exchange standards and metadata descriptions common to all service providers are fundamental keys to success. On a technical point of view, XML may be one important tool for data integration. {\\sc Aladin} is an example of a comprehensive tool, allowing integration of reference images, with information from catalogues, databases and data archives.  The rapid development of the world-wide `bibliographic network' has been particularly impressive, and the `data archive network' seems well under way. The CDS is a major hub in the on-line `Virtual Observatory' presently under construction: its services allow astronomers to select the information of interest for their research, and to access original data, observatory archives and results published in journals."
    },
    "0002/gr-qc0002089_arXiv.txt": {
        "abstract": "We study the evolution of gravitational waves through the preheating era that follows inflation. The oscillating inflaton drives parametric resonant growth of scalar field fluctuations, and although super-Hubble tensor modes are not strongly amplified, they do carry an imprint of preheating. This is clearly seen in the Weyl tensor, which provides a covariant description of gravitational waves. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002231_arXiv.txt": {
        "abstract": "Two ways of possibly solving the angular momentum problem plaguing cold dark matter (CDM) {\\it ab initio} simulations of disk galaxy formation are discussed: 1) Stellar feedback processes and 2) Warm dark matter (WDM) rather than CDM.  In relation to the chemical evolution of disk galaxies our simulations indicate that in case 1) the first generation of {\\it disk} stars formed in disk galaxies like the Milky Way should have an abundance about two dex below solar, in fairly good agreement with the lowest observed abundance of the metal-weak tail of the Galactic thick disk. For the second case no such statements can be made without further assumptions about the star-formation history of the galaxies.  We find that the $I$-band Tully-Fisher relation can be matched by WDM disk galaxy formation simulations provided $(M/L_I) \\sim$ 0.8 for disk galaxies, which Sommer-Larsen \\& Dolgov (1999) argue is a reasonable value.  Finally it is discussed how the magnetic field strengths observed in galactic disks can be obtained through disk galaxy {\\it formation}, as an alternative to the conventional dynamo hypothesis. ",
        "introduction": "\\label{s:intro}  The formation of galactic disks is one of the most important unsolved problems in astrophysics today. In the currently favored hierarchical clustering framework, disks form in the potential wells of dark matter halos as the baryonic material cools and collapses dissipatively. Fall \\& Efstathiou (1980) have shown that disks formed in this way can be expected to possess the observed amount of angular momentum (and therefore the observed spatial extent for a given mass and profile shape), but only under the condition that the infalling gas retain most of its original angular momentum.  However, numerical simulations of this collapse scenario in the cold dark matter (CDM) cosmological context (e.g., Navarro \\& Benz 1991, Navarro \\& White 1994, Navarro, Frenk, \\& White 1995) have so far consistently indicated that when only cooling processes are included the infalling gas loses too much angular momentum (by over an order of magnitude) and the resulting disks are accordingly much smaller than required by the observations. This discrepancy is known as the {\\em angular momentum problem} of disk galaxy formation. It arises from the combination of the following two facts: a) In the CDM scenario the magnitude of linear density fluctuations $\\sigma(M) = \\langle(\\delta M/M)^2\\rangle^{1/2}$ increases steadily with decreasing mass scale $M$ leading to the formation of non-linear, virialized structures at increasingly early epochs with decreasing mass i.e.~the hierarchical ``bottom-up'' scenario.  b) Gas cooling is very efficient at early times due to gas densities being generally higher at high redshift as well as the rate of inverse Compton cooling also increasing very rapidly with redshift.  a) and b) together lead to rapid condensation of small, dense gas clouds, which subsequently lose energy and (orbital) angular momentum by dynamical friction against the surrounding dark matter halo before they eventually merge to form the central disk. A mechanism is therefore needed that prevents, or at least delays, the collapse of protogalactic gas clouds and allows the gas to preserve a larger fraction of its angular momentum as it settles into the disk. Two such possible solutions are discussed in section 2.  In section 3 we present some new results from our WDM disk galaxy formation simulations on the Tully-Fisher relation and in section 4 we discuss how the magnetic field strengths of a few $\\mu$G observed in galactic disks can be obtained via disk galaxy {\\it formation}, as an alternative to disk dynamo amplification. ",
        "conclusions": ""
    },
    "0002/astro-ph0002141_arXiv.txt": {
        "abstract": "The detection and first identification of the deuterium Balmer emission lines, \\da\\ and \\db, in the core of the Orion Nebula is reported. These lines are very narrow, have identical 11~\\kms\\ velocity shifts with respect to \\ha\\ and \\hb, are probably excited by UV continuum fluorescence from the Lyman (\\ion{D}{i}) lines and arise from the interface between the \\hii\\ region and the molecular cloud. ",
        "introduction": "Deuterium is believed to be entirely produced in the Big Bang and then steadily destroyed by astration (Epstein et al.~\\cite{epstein76}). Standard models predict a decrease of its abundance by a factor 2--3 in 15~Gyrs (\\eg,~Tosi et al.~\\cite{tosi98}). This picture is essentially constrained by deuterium abundance determinations at $\\sim15$~Gyrs (primordial intergalactic clouds), 4.5~Gyrs (protosolar) and 0.0~Gyrs (interstellar medium). Although the evolution of the deuterium abundance seems to be qualitatively understood, the measurements show some dispersion. Thus, absorption in the Lyman series provides interstellar deuterium abundance \\dshism$\\simeq1.5\\times10^{-5}$ (Linsky~\\cite{linsky98}), but with fluctuations that may well be real (Vidal-Madjar et al. 1998). These dispersions led to the development of non-standard models in which, for example, deuterium may either decrease by more than a factor 4 in 15~Gyrs (\\eg,~Vangioni-Flam et al.~\\cite{flam94}) or be created/destroyed by new mechanisms [\\eg, Lemoine et al.~(\\cite{lemoine99}) for a~review].  A detailed appraisal of the evolution of deuterium is crucial for cosmology and galactic chemical evolution. The most reliable estimate of \\dshism\\ to date is based on far-UV observation from space (Copernicus, IMAPS, HST or FUSE) of the Lyman lines of D and H in absorption. These lines are also observed in the optical and near-UV to obtain D/H in high redshift quasar absorbers. Other D/H determinations include {\\it in situ} measurements in the Solar System (\\eg,~Mahaffy et al.~\\cite{mahaffy98}), observations of molecules such as HD or DCN (\\eg,~Bertoldi et al.~\\cite{bertoldi99}) and observations of \\ion{D}{i}~92~cm (\\eg,~Chengalur et al.~\\cite{chengalur97}).  New methods to determine D/H are of interest. One possibility is ground-based observation of the deuterium lines. The isotope shift of the deuterium Balmer lines with respect to the hydrogen Balmer lines is $-81.6$~\\kms. These \\di\\ lines have never been identified before. Attempts to detect \\da\\ in absorption in the Sun (Beckers~\\cite{beckers75}) and early-type stars (\\eg,~Vidal-Madjar et al.~\\cite{avm88}) were unsuccessful (D is destroyed in stars). Traub et al.~(\\cite{traub74}) observed \\ha\\ in the Orion Nebula using three-etalon Fabry-Perot spectrometers and reported D/H upper limits.  Here we report on spectra of Orion, secured at the Canada-France-Hawaii Telescope (CFHT). Emission lines detected in the blue wings of \\ha\\ and \\hb\\ are identified with \\da\\ and \\db. A preliminary account was presented by H\\'ebrard et al. (\\cite{hebrard99}). Observations are described in Sect.~\\ref{observations}, the identification and the origin of the lines in Sect.~\\ref{identification} and \\ref{origin} and the excitation mechanism in Sect.~\\ref{fluo}. ",
        "conclusions": "Deuterium is identified for the first time in a nebula from optical spectroscopy. The excitation mechanism of the observed lines, \\da\\ and \\db, is continuum fluorescence from Ly$_{\\rm D}$ lines in the PDR. Considering the saturation of the first Ly$_{\\rm D}$ lines and the possible influence of the neighbouring Ly$_{\\rm H}$ lines, observing the full deuterium Balmer series is essential to obtain a D/H value from optical data and appears feasible, at least in Orion.  An optical determination of D/H in \\hii\\ regions would allow to check existing \\dshism\\ values and obtain D/H in low-metallicity extragalactic \\hii\\ regions, where the deuterium abundance should be close to its primordial~value.  The large photoexcitation cross section of the first Lyman lines makes deuterium Balmer fluorescence a sensitive way to {\\sl detect} deuterium in nebulae, leading for example to stringent upper limits to D/H in planetary nebulae (H\\'ebrard et al.~\\cite{hebrard00b}), where deuterium is depleted."
    },
    "0002/astro-ph0002188_arXiv.txt": {
        "abstract": "The basic workings of inflationary models are summarized, along with the arguments that strongly suggest that our universe is the product of inflation.  The mechanisms that lead to eternal inflation in both new and chaotic models are described.  Although the infinity of pocket universes produced by eternal inflation are unobservable, it is argued that eternal inflation has real consequences in terms of the way that predictions are extracted from theoretical models.  The ambiguities in defining probabilities in eternally inflating spacetimes are reviewed, with emphasis on the youngness paradox that results from a synchronous gauge regularization technique.  To clarify (but not resolve) this ambiguity, a toy model of an eternally inflating universe is introduced.  Vilenkin's proposal for avoiding these problems is also discussed, as is the question of whether it is meaningful to discuss probabilities for unrepeatable measurements. ",
        "introduction": "\\setcounter{equation}{0}  There are many fascinating issues associated with eternal inflation, which will be the main subject of this talk. You have certainly heard other people talk about eternal inflation, but I feel that the topic is important enough so that you should hear about it in some accent other than Russian.  I will begin by summarizing the basics of inflation, including a discussion of how inflation works, and why many of us believe that our universe almost certainly evolved through some form of inflation.  This material is certainly not new, but I think it is an appropriate introduction to any volume that focuses on inflationary cosmology.  Then I will move on to discuss eternal inflation, first explaining how it works.  I will then argue the eternal inflation has important implications, and raises important questions, which should not be dismissed as being merely metaphysical. ",
        "conclusions": "\\setcounter{equation}{0}  In this paper I have summarized the workings of inflation, and the arguments that strongly suggest that our universe is the product of inflation.  I argued that inflation can explain the size, the Hubble expansion, the homogeneity, the isotropy, and the flatness of our universe, as well as the  absence of magnetic monopoles, and even the characteristics of the nonuniformities. The detailed observations of the cosmic background radiation anisotropies continue to fall in line with inflationary expectations, and the evidence for an accelerating universe fits well with the inflationary preference for a flat universe.  Next I turned to the question of eternal inflation, claiming that essentially all inflationary models are eternal. In my opinion this makes inflation very robust: if it starts anywhere, at any time in all of eternity, it produces an infinite number of pocket universes.  Eternal inflation has the very attractive feature, from my point of view, that it offers the possibility of allowing unique predictions even if the underlying string theory does not have a unique vacuum.  I have also emphasized, however, that there are important problems in understanding the implications of eternal inflation.  First, there is the problem that we do not know how to treat the situation in which the scalar field climbs upward to the Planck energy scale.  Second, the definition of probabilities in an eternally inflating spacetime is not yet a closed issue, although important progress has been made.  And third, I might add that the entire present approach is at best semiclassical.  A better treatment may not be possible until we have a much better handle on quantum gravity, but eventually this issue will have to be faced."
    },
    "0002/astro-ph0002377_arXiv.txt": {
        "abstract": "We estimate the quantum corrections to the ground state energy in neutron matter (which could be termed as well either shell correction energy or Casimir energy) at subnuclear densities, where various types of inhomogeneities (bubbles, rods, plates) are energetically favorable. We show that the magnitude of these energy corrections are comparable to the energy differences between various types of inhomogeneous phases. We discuss the dependence of these corrections on a number of physical parameters (density, filling factor, temperature, lattice distortions). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002411_arXiv.txt": {
        "abstract": "A kinetic model of particle acceleration in supernova remnants (SNRs) is extended to study the cosmic ray (CR) and associated high-energy $\\gamma$-ray production during SN shock propagation through the inhomogeneous circumstellar medium of a progenitor star that emits a wind. The wind forms a low-density bubble surrounded by a swept-up shell of interstellar matter. $\\gamma$-rays are produced as a result of decay of pions which in turn are the result of collisions of CRs with nuclei of the thermal plasma. The time evolution of the SNRs is followed numerically, taking into account the nonlinear backreaction of the accelerated CRs. The model for SNRs includes injection of suprathermal particles at the shock front and heating of the thermal plasma due to dissipation of Alfv\\'en waves in the precursor region. Examples typical for SN type Ib and SN type II explosions are considered. Apart from the confirmation of the known fact that acceleration is extremely rapid and that the upper momentum cutoff is reached almost immediately after the explosion due to the high wind magnetic field, it is also shown that the CRs are accelerated with a high efficiency. Depending on the circumstellar parameters, 20\\% to 40\\% of the SN explosion energy is absorbed by CRs during the SNR evolution for a moderate injection rate, when a fraction $\\eta = 10^{-3}$ of the gas particles, swept up by the supernova shock, is accelerated. The CR momentum spectrum, ultimately produced in the SNRs, has a power law form $N\\propto p^{-\\gamma}$ with an index $\\gamma=2.0$ to 2.1 in a wide energy range up to a maximum energy, which is about $10^{14}$~eV, if the CR diffusion coefficient is as small as the Bohm limiting value. It is to be expected that the resulting CR chemical composition at high energies reflects more the stellar wind composition, whereas at lower energies it corresponds more to that of the average interstellar medium. The expected $\\pi^0$-decay $\\gamma$-ray flux is however considerably lower than in the case of a uniform interstellar medium; a relatively high $\\gamma$-ray luminosity in the band $\\epsilon_{\\gamma}\\gsim 1$~TeV, detectable at distances of several kpc, is only expected in the case of a relatively dense ISM with a number density above 10~cm$^{-3}$.  Extremely high $\\gamma$-ray emission may be produced when the SN shock propagates through the slow dense wind of a red supergiant, the progenitor of a SN type II. In this case SNRs might be visible in $\\gamma$-rays for several hundred years out to distances of tens of kpc. For the case of a SN type Ib the expected $\\pi^0$-decay {\\gr} TeV-energy flux during the whole SNR evolution remains lower than $10^{-11}$~cm$^{-2}$s$^{-1}$ if the interstellar number density is less than $0.3$~cm$^{-3}$. ",
        "introduction": "Considerable efforts have been made during the last years to empirically confirm the theoretical expectation that the main part of the Galactic cosmic rays (CRs) originates in supernova remnants (SNRs). Theoretically progress in the solution of this problem has been due to the development of the theory of diffusive shock acceleration (see, for example, reviews by Drury 1983; Blandford \\& Eichler 1987; Berezhko \\& Krymsky 1988). Although still incomplete, the theory is able to explain the main characteristics of the observed CR spectrum under several reasonable assumptions, at least up to an energy of $10^{14}\\div 10^{15}$~eV. Direct information about the dominant nucleonic CR component in SNRs can only be vobtained from $\\gamma$-ray observations. If this nuclear component is strongly enhanced inside SNRs then through inelastic nuclear collisions, leading to pion production and subsequent decay, $\\gamma$-rays will be produced.  CR acceleration in SNRs expanding in a uniform interstellar medium (ISM) (Drury et al. 1989; Markiewicz et al. 1990; Dorfi 1990), and the properties of the associated $\\gamma$-ray emission (Dorfi 1991; Drury et al. 1994) were investigated in a number of studies (we mention here only those papers which include the effects of shock geometry and time-dependent nonlinear CR backreaction; for a review of others which deal with the test particle approximation, see for example Drury 1983; Blandford \\& Eichler 1987; Berezhko \\& Krymsky 1988; V\\\"olk 1997). All of these studies are based on a two-fluid hydrodynamical approach and directly employ the assumption that the expanding SN shock is locally plane; as dynamic variables for the CRs the pressure and the energy density are determined. Their characteristics are sometimes essentially different from the results obtained in a kinetic approach (Kang \\& Jones 1991; Berezhko et al. 1994, 1995, 1996) which consistently takes the role of shock geometry and nonlinear CR backreaction into account. First of all, in kinetic theory the form of the spectrum of accelerated CRs and their maximum energy are calculated selfconsistently. In particular, the maximum particle energy $\\epsilon_{max}$, achieved at any given evolutionary stage, is determined by geometrical factors (Berezhko 1996), in contrast to the hydrodynamic models which in fact postulate that the value of $\\epsilon_{max}(t)$ is determined by the time interval $t$ that has passed since the explosion (Drury et al. 1989; Markiewicz et al. 1990; Dorfi 1990; Jones \\& Kang 1992). Although the difference between the values of $\\epsilon_{max}$ in the two cases is not very large, it critically influences the structure and evolution of the shock. For example, the shock never becomes completely modified (smoothed) by the CR backreaction (Kang \\& Jones 1991; Berezhko et al. 1994, 1995, 1996). Together with the smooth precursor, the shock transition always contains a relatively strong subshock which heats the swept-up gas and leads to the injection of suprathermal gas particles into the acceleration process. In this sense diffusive shock acceleration is somewhat less efficient than predicted by hydrodynamic models. Acceleration always requires some freshly injected particles which are generated during gas heating. This prediction is in agreement with the observations that show significant gas heating in young SNRs.  On the other hand, the shock modification by the CR backreaction is even greater than predicted by hydrodynamic models. The total shock compression ratio $\\sigma \\propto M^{3/4}$ (Berezhko et al. 1996) is a monotonically rising function of the shock Mach number $M$ and can significantly exceed the classical value 4, or even the value $\\sigma =7$ that corresponds to a postshock medium dominated by relativistic CRs. This result is a direct consequence of the fact that CR acceleration at the front of a spherically expanding shock is accompanied by a temporally increasing dilution of shock energy in the form of CRs in the precursor region (Drury et al. 1995; Berezhko 1996). Qualitatively this leads to the same effect as if energetic particles or thermal gas energy were radiated away from the shock (Berezhko 1996; Berezhko \\& Ellison 1999). We note that this energy dilution effect is much stronger than the effect of a decreasing specific heat ratio due to the conversion of shock energy into relativistic CRs.  The additional gas heating in the precursor region due to Alfv\\'en wave dissipation significantly restricts the shock compression ratio $\\sigma(t)$ even though its value still considerably exceeds 4 for a strong shock (Berezhko et al. 1996; Berezhko \\& V\\\"olk 1997). The precursor gas heating has a much less pronounced effect in the hydrodynamic description (e.g. Dorfi 1990).  As far as the expected $\\gamma$-ray emission, produced in SNRs by the nuclear CR component, is concerned, there are less significant differences between the kinetic (Berezhko \\& V\\\"olk 1997) and the hydrodynamic (Dorfi 1991; Drury et al. 1994) predictions, even though these differences are not unimportant. The reason is that the peak value of the $\\gamma$-ray flux is mainly determined by the fraction of overall hydrodynamic explosion energy that is transferred to CRs, i.e. by the overall efficiency of CR acceleration which is not strongly dependent on the model used.  The differences in the time variation of the predicted $\\gamma$-ray fluxes are more essential. Kinetic theory (Berezhko \\& V\\\"olk 1997) revealed much more effective CR and $\\gamma$-ray production during the free expansion phase. It also shows a more rapid decrease of the $\\gamma$-ray flux during the subsequent Sedov phase after reaching its peak value, due to the different spatial distributions of thermal gas and CRs inside SNRs. This energy-dependent lack of overlap is not taken into account in hydrodynamic models.  The model developed by Ellison and co-workers to describe diffusive shock acceleration of CRs and to predict the expected $\\gamma$-ray emission from SNRs, uses a Monte Carlo (MC) simulation of pitch angle scattering. It is applied in the same way to all particles, both CR and thermal particles, and is correspondingly solved numerically from the outset (e.g. Ellison et al. 1995; Baring et al. 1999). The energetic particle population (i.e. accelerated CRs) naturally arises in this approach as a high energy tail of the distribution of gas particles which undergo heating in the strongest gradient of the selfconsistently modified shock transition, the subshock in our notation. In this way the model incorporates not only selfconsistent CR acceleration but also gas heating and particle injection into the acceleration process. Since it is a plane-wave steady state model it includes as a free parameter the value of the CR cutoff energy $\\epsilon_{max}$, which can not be calculated in this kind of model. On the other hand, at sufficiently high particle energies the description corresponds to a diffusive approach based on the diffusive transport equation. Therefore our time-dependent kinetic description in spherical symmetry, with an injection rate that corresponds to the MC-model, and the MC-model, in turn incorporating the same value of CR cutoff energy, give identical CR spectra at any stage of the free expansion phase for the same set of ISM and shock parameters (Berezhko \\& Ellison 2000). This means that in its free expansion phase the SNR evolution can be represented as a sequence of quasi-stationary states; each of them can be reproduced in the framework of the plane-wave steady state description with the corresponding value of the CR cutoff energy $\\epsilon_{max}(t)$. The situation becomes more complicated in the subsequent Sedov phase due to the existence of a so-called escaping particle population near the shock front (Berezhko \\& Ellison 1999b), which is a purely nonstationary phenomenon (Berezhko 1986a; Berezhko \\& Krymsky 1988) and therefore can not be reproduced in the framework of a steady state approach. As far as the overall CR spectrum and the associated $\\gamma$-ray flux at any given evolutionary phase are concerned, the plane wave description can only give an approximate estimate because it does not include such important physical factors as CR adiabatic cooling and the incomplete overlap of the spatial distributions of gas and CRs inside the expanding SNR.  The evolution of SNRs in a uniform ISM is typical only for SNe type Ia. SNe of type Ib and II, which are more numerous in our Galaxy, explode into an inhomogeneous circumstellar me\\-di\\-um, formed by the intensive wind of their massive progenitor stars (Losinskaya 1991). This is at least true for progenitor masses $M \\gsim 20 M_{\\odot}$ whose main sequence winds are very energetic. Their time-integrated hydrodynamic energy output is so large that they create wind cavities which are about as large as the size of a remnant of the subsequent Supernova explosion, if it would occur in a uniform ISM of the same density. The more frequent lower mass core collapse progenitors ($8 \\lsim M \\l 20 M_{\\odot}$) have at least a massive Red Supergiant wind region before the transition to the external ISM. Therefore the description of SNR evolution and CR acceleration becomes considerably more difficult for those cases. Due to the complicated structure of the medi\\-um one can expect that the difference between kinetic and hydrodynamic model predictions will be even more pronounced than in the case of a uniform ISM.  Up to now, except for the paper by Berezhko \\& V\\\"olk (1995) where preliminary results of the present work were presented, the consideration of CR acceleration in the case of a SN type Ib and SN type II was considered only for SN shock propagation through the region of the supersonic stellar wind (Berezinskii \\& Ptuskin 1988; V\\\"olk \\& Biermann 1988; Jones \\& Kang 1992; Kirk et al. 1995). Since this region is characterized by an initially very strong if spatially decreasing magnetic field of stellar origin and a very high shock speed, the maximum energy of CRs accelerated in the wind region is achieved extremely rapidly during a period of time comparable with the SNR age at each phase, and then it remains nearly constant due to a balance of acceleration gains and adiabatic losses (V\\\"olk \\& Biermann 1988). As it is shown below, the expected CR energy spectrum is close to being selfsimilar where the amplitude monotonically decreases with time but the shape remains invariant (see Appendix). Thus the situation is very different from that for a uniform circumstellar medium with a standard magnitude ISM field strength, where the maximum CR energy is only achieved at the end of the sweep-up phase, after hundreds or even thousands of years.  The direct dependence on the stellar surface field also raised hopes that the CR acceleration in the wind would constitute a major source of the observed Galactic CRs with energies $\\epsilon \\gsim 10^{14}$~eV. The most optimistic assumptions predict CR generation in these regions up to an energy of $3\\times 10^{18}$~eV (Biermann 1993). In this context we would like to make the following remarks:  First of all, the supersonic wind region contains a relatively small amount of mass $M\\lsim 1 M_{\\odot}$ which is typically smaller than the ejected mass $M_{ej}\\sim 10M_{\\odot}$. Therefore the fraction of the SN explosion energy which can be transfered into CRs is typically lower than required for the observed CR spectrum, especially in the case of a SN type Ib (see below).  Second, to achieve a particle energy significantly greater than $10^{14}$~eV one needs to assume an unusually high magnetic field in the wind. This tends to violate the condition $V_{w} \\geq c_{a}$, necessary for the existence of a supersonic flow in the first place, as pointed out by Axford (1994), where $V_w$ is the wind speed and $c_a$ is the Alfv\\'en velocity.  Third, the wind magnetic field is largely azimuthal, except in the very polar region near the stellar rotation axis, and the outer SNR shock is almost perpendicular. Therefore not only the suprathermal ion injection is much less efficient but also the CR acceleration itself appears not so efficient as in the case of a quasi-parallel shock. Although one can give some physical arguments which assert that particle injection with their subsequent acceleration takes place also in this case, one needs a more rigorous treatment of this complex question. An example for the acceleration at a shock propagating through a stationary stochastic large yyscale magnetic field that is on average perpendicular to the shock normal is given by Kirk et al. (1996) based on the process of anomalous diffusion (see Chuvilgin \\& Ptuskin 1993). This test particle calculation indicates, as intuitively expected, that the spectrum of accelerated particles is steeper than for acceleration at a quasi-parallel shock. As a consequence, the nonlinear efficiency of acceleration would also have to be assumed to be significantly reduced.  Fourth, the unusually high maximum CR energy predicted by Biermann (1993) is a direct consequence of the assumption of a very small and energy independent CR diffusion coefficient, presumably produced by the chaotic turbulent motions of the medium near the shock, that formally gives very fast acceleration to extremely high energies. Leaving aside the important question concerning the justification of the assumption for such shock-produced turbulent motions also significantly ahead of the forward SNR shock itself (see e.g. Ellison et al. 1994, Lucek \\& Bell 2000), one has to stress that the final result should in any case depend upon the values of the microscopic CR diffusion tensor which alone allows CRs to intersect the shock front many times with a subsequent increase of their energy. As far as the macroscopic CR diffusion due to chaotic gas motions (Biermann 1993) is concerned, it can not itself provide CR acceleration. In fact, from what has been said before, even our results presented below, based on the usual assumption about Bohm type microscopic particle diffusion in a quasi-parallel shock and much less optimistic compared with those in Biermann (1993), must be considered as an upper limit for the CR energy density which can be achieved during shock acceleration. This leaves open the question whether the evolving SNR shock system can generate a strong scattering wave field and allows the assumption of a favorable ratio of the microscopic parallel and perpendicular diffusion coefficients over a sufficiently long time, so that the maximum individual particle energy can be increased by something like an order of magnitude or more for a perpendicular shock (Jokipii 1987; Ostrowski 1988; Ellison et al. 1995; Reynolds 1998).  As mentioned above, a strong wind modifies the environment of massive progenitors of SN types Ib and II. It sweeps up the ambient gas into a thin shell surrounding a rarefied bubble (Weaver et al. 1977; Losinskaya 1991). Since the typical cavity is greater than $10$~pc in size and contains a considerable amount of matter, it can significantly influence the SNR evolution and CR acceleration. Estimates and preliminary numerical calculations indicate that the SN shock propagation through the progenitor's wind region should generate high-energy $\\pi^0$-decay $\\gamma$-ray emission on a detectable level (Berezinsky \\& Ptuskin 1988; Kirk et al. 1995; Berezhko \\& V\\\"olk 1995\\footnote{Due to an error in the plot routine in Fig. 1 of Berezhko \\& V\\\"olk (1995), the \\gr fluxes in the wind SNR cases were plotted with an enlargement factor of 2500 which should be left out}). Yet, up to now, all {\\it detailed} numerical investigations of CR acceleration in SNRs have dealt with the case of a star exploding into a uniform homogeneous medium.  In this paper we present a detailed extension of the kinetic model for CR acceleration in SNRs (Berezhko et al. 1994, 1995, 1996; Berezhko \\& V\\\"olk 1997) to the case of a nonuniform circumstellar medium, spherical symmetry still being assumed. We study the CR acceleration and $\\pi^0$-decay $\\gamma$-ray production taking into account that the circumstellar medium can be strongly modified by a wind from the progenitor star. In doing this we assume energetic particle scattering governed by the Bohm diffusion coefficient in the local mean magnetic field, treating the shock normal to be quasi-parallel to the field direction. We have studied two cases typical for type Ib and type II SNe. In the first case, a star with an initial mass of $35~M_{\\odot}$ at the end of its evolution explodes as SN into the cavity created by a main-sequence (MS) O-star wind and the subsequent winds during the red supergiant (RSG) and Wolf-Rayet (WR) phases. For the case of a SN type II we take the example of the explosion of a star with initial mass $15~M_{\\odot}$ into the cavity created by the winds emitted during the MS and RSG phases. In both cases we consider the two different interstellar number densities 0.3 and 30 cm$^{-3}$.  We shall not address here the question of $\\gamma$-ray production by electron synchrotron emission, Bremsstrahlung, or the inverse Compton (IC) effect on ambient low energy photons. Especially at very high energies exceeding $\\sim 10$~GeV, IC $\\gamma$-rays appear to dominate the emission from plerions like the Crab nebula (e.g. de Jager \\& Harding 1992, Atoyan \\& Aharonian 1996), where presumably a relativistic wind of electron-positron pairs from the pulsar is dissipated in a circumstellar termination shock deep inside the SNR shell. The IC effect associated with acceleration of ultrarelativistic electrons inside the SNR shell, or at its leading shock (along with ions considered here), may also contribute quantitatively to the $\\gamma$-ray luminosity of a SNR as simple estimates suggest (Mastichiadis 1996; Mastichiadis \\& de Jager 1996). Therefore in any specific source this leptonic contribution to the $\\gamma$-ray flux needs to be estimated before the hadronic $\\pi^0$-decay $\\gamma$-ray emission can be compared with the theoretical models presented below.  We briefly describe some aspects of the model in Sect. 2 since it was in detail described in a previous paper (Berezhko \\& V\\\"olk 1997). Sect. 3 contains the results and Sect. 4 includes the discussion and the conclusions. ",
        "conclusions": "Our numerical results show that when a SN explodes into a circumstellar medium strongly modified by a wind from a massive progenitor star, then CRs are accelerated in the SNR almost as effectively as in the case of a uniform ISM (Berezhko et al. 1994, 1995, 1996; Berezhko \\& V\\\"olk 1997): about $20\\div 40$\\% of the SN explosion energy is transformed into CRs during the active SNR evolution.  During SN shock propagation in the supersonic wind region very soon the acceleration process reaches a quasi-stationary level which is characterized by a high efficiency and a correspondingly large shock modification. Despite the fact that the shock modification is much stronger than predicted by a two-fluid hydrodynamical model (Jones \\& Kang 1992), the shock never becomes completely smoothed by CR backreaction: a relatively strong subshock always exists and plays an important dynamical role. As in the case of a uniform ISM, the spurious complete shock smoothing in hydrodynamic models is the result of an underestimate of the role of geometric factors.  Due to the relatively small mass contained in the supersonic wind region CRs absorb there only a small fraction of the explosion energy (about 1\\% in the case of a SN type Ib, and 10\\% in the case of a SN type II) and the SNR is still very far from the Sedov phase after having swept up this region. Therefore we conclude, that the CRs produced in this region should not play a very significant role for the formation of the observed Galactic CR energy spectrum.  The peak value of the CR energy content in the SNR is reached when the SN shock sweeps up an amount of mass roughly equal to several times the ejected mass. This takes place during the SN shock propagation in the modified bubble. In a purely adiabatic bubble the sweep-up would occur at the beginning of the shell. Compared with the uniform ISM case the subsequent adiabatic CR deceleration is less important in the case of a modified circumstellar medium. The main amount of CRs in this case is produced when the SN shock propagates through the bubble. In this stage the dynamical scale length is much smaller than the shock size. Therefore the relative increase of the shock radius during the late evolution stage and the corresponding adiabatic effects are small. These configurational properties lead to potentially interesting changes of the CR chemical composition with particle energy (see Sect. 3.1)  In the case of the modified circumstellar medium the CR and $\\gamma$-ray spectra are more variable during the SN shock evolution than in the case of a uniform ISM (Berezhko \\& V\\\"olk 1997). At the same time the form of the resulting overall CR spectrum is rather insensitive to the parameters of the ISM as in the case of uniform ISM.  The reason is that the main amount of CRs are produced in the latest phase which has still a strong enough shock. Roughly speaking, the overall CR spectrum (except the most energetic CRs) is mainly formed at the stage when the shock compression ratio lies between 4 and 5.  The maximum energy of the accelerated CRs reached during the SNR evolution is about $10^{14}$~eV for protons in all the cases considered.  Our results confirm the important conclusion, reached for the case of a uniform ISM before, that the diffusive acceleration of CRs in SNRs is able to generate the observed CR spectrum up to an energy $\\sim 10^{14}$~eV, if the CR diffusion coefficient is as small as Bohm limit. This disregards the possibility of turbulent field amplification, as discussed in Sect. 2.  In the case of a SN Ib the expected TeV-energy $\\gamma$-ray flux, normalized to a distance of 1 kpc, remains lower than $10^{-12}$~cm$^{-2}$s$^{-1}$ during the entire SNR evolution if the ISM number density is less than 1~cm$^{-3}$ except for an initial short period $t<100$~yr when it is about $10^{-11}$~cm$^{-2}$s$^{-1}$.  Only for a relatively dense ISM with $N_H=30$~cm$^{-3}$ the expected $\\gamma$-ray flux is about $10^{-10}$~cm$^{-2}$s$^{-1}$ at late phases $t>10^4$~yr. A similar situation exists at late phases of SNR evolution in the case of SN II. It is interesting to note that the expected $\\gamma$-ray flux is considerably lower, at least by a factor of hundred, compared with the case of uniform ISM of the same density $N_H$. This confirms the preliminary result reported earlier (Berezhko \\& V\\\"olk 1995)  The type II SN explodes into the dense wind of the red supergiant progenitor star. During the first several hundred years $t_m$ after the explosion, the expected TeV-energy $\\gamma$-ray flux at a distance $d=1$~kpc exceeds the value $10^{-9}$~cm$^{-2}$s$^{-1}$ and can be detected up to the distance $d_m=30$~kpc with present instruments like HEGRA, Whipple or CAT. This distance is of the order of the diameter of the Galactic disk (see also Kirk et al. 1995). Therefore all Galactic SNRs of this type whose number is $N_{sn}=\\nu_{sn}t_m$ should be visible. But in this case we can expect at best $N_{sn}\\sim 10$ such $\\gamma$-ray sources at any given time.  The typical value of the cutoff energy of the expected $\\gamma$-ray flux is about $10^{13}$~eV, if the CR diffusion coefficient is as small as the Bohm limit. In this respect the negative result of high-threshold arrays (Borione et al. 1995; Allen, G.E. et al. 1995; Allen, N.H. et al. 1995) in searching of $\\gamma$-ray emission from Galactic SNRs is not surprising because their threshold $E_{th}\\sim 50$~TeV exceeds the cutoff energy of the expected $\\gamma$-ray flux; marginally this also holds for the negative results of the lower threshold $E_{th}>$~20 TeV AIROBICC array (Prahl \\& Prosch 1997; Prosch et al. 1996). It is less obvious how to interpret the negative results of imaging atmospheric Cherenkov telescopes with thresholds less than about 1 TeV (Mori et al. 1995; Lessard et al. 1997; Hess et al. 1997; Hess 1998; Buckley et al. 1998). For core collapse SN of types II or Ib with quite massive progenitors one can in part explain this fact by the extremely low $\\pi^0$-decay $\\gamma$-ray intensity expected from such SNRs during the period of SN shock propagation through the low-density hot bubble. An alternative possibility relates to the assumption of the Bohm limit for the CR diffusion coefficient which can be too optimistic, in particular for the quai-perpendicular geometry in wind-blown regions from rotating stars. For a slightly more general discussion of SNR $\\gamma$-rays in stellar wind cavities, see V\\\"olk (1997)."
    },
    "0002/astro-ph0002282_arXiv.txt": {
        "abstract": "We present mid--infrared spectro--imagery and high--resolution spectroscopy of the Orion bar and of a region in the Orion nebula. These observations have been obtained in the Guaranteed Time with the Circular Variable Filters of the ISO camera (CAM-CVF) and with the Short Wavelength Spectrometer (SWS), on board the European Infrared Space Observatory (ISO). Our data shows emission from amorphous silicate grains from the entire \\HII~ region and around the isolated O9.5V star $\\theta^2$~Ori~A. The observed spectra can be reproduced by a mixture of interstellar silicate and carbon grains heated by the radiation of the hot stars present in the region. Crystalline silicates are also observed in the Orion nebula and suspected around $\\theta^2$~Ori~A. They are probably of interstellar origin.  The ionization structure and the distribution of the carriers of the Aromatic Infrared Bands (AIBs) are briefly discussed on the basis of the ISO observations. \\keywords\t{ISM: Orion nebula\t\t- ISM: Orion bar  \t\t- ISM: \\HII~ regions\t\t- stars: $\\theta^2$~Ori~A\t\t- dust, extinction\t\t- Infrared: ISM: lines and bands} ",
        "introduction": "The Orion nebula is one of the most studied star--forming regions in the Galaxy.  The ionizing stars of the Orion nebula (the Trapezium stars, the hottest of which is $\\theta^1$~Ori~C, O6) have eroded a bowl--shaped \\HII~ region into the surface of the Orion molecular cloud. The Orion bar is the limb--brightened edge of this bowl where an ionization front is progressing into the molecular cloud. It is seen as an elongated structure at a position angle of approximately 60\\degr. The Orion nebula extends to the North.  The Trapezium stars are located at an angular distance of approximately 2.3 arc minutes from the bar, corresponding to 0.35 pc at a distance of 500 pc.  The molecular cloud extends to the other side of the bar, but also to the back of the Orion nebula. The bright star $\\theta^2$~Ori~A (O9.5Vpe) lies near the bar, and is clearly in front of the molecular cloud since its color excess is only E(B--V) $\\simeq$ 0.2 mag.  \\begin{figure*}[t!] {\\psfig{file=orion_minimosaic.ps,width=17cm,angle=00}} \\caption{Mosaic of six images of the Orion bar area (shown in detector coordinates, a clockwise rotation of 10.4\\degr~ is needed to display the real sky orientation; see Fig. 2 for the equatorial coordinates). Top row: {\\it (left)} an image at 5.01\\mum, $\\theta^2$~Ori~A is visible in the middle left of the image; {\\it (centre)} the Orion bar at the AIB wavelength of 6.2\\mum; {\\it (right)} image at 9.5\\mum (the wavelength of one of the silicate features), note that $\\theta^2$~Ori~A is again visible. Bottom row: {\\it (left)} image at the AIB wavelength of 11.3\\mum~ and {\\it (centre)} image at 12.7\\mum~(ISOCAM's CVF resolution blends [\\NeII] and the 12.7\\mum~AIB feature); image at 15.6\\mum {\\it (right)}, wavelength of the [\\NeIII] forbidden line.  } \\end{figure*}  Figs. 1 and 2 illustrate the geometry of the region observed. Fig. 1 shows six representative images of the region of the Orion bar (see the figure caption for details). Fig. 2 shows the contours of the the [\\NeIII] 15.5\\mum~ fine-structure line emission which delineates the \\HII~ region. The emission in one of the mid-IR bands (hereafter called the Aromatic Infrared Bands, AIBs) at 6.2\\mum~ traces the Orion bar (an edge-on PhotoDissociation Region or PDR). The AIBs are usually strongly emitted by PDRs.  The Trapezium region was avoided because of possible detector saturation.  Pioneering infrared observations by Stein \\& Gillet (\\cite{Stein}) and Ney \\etal (\\cite{Ney}) discovered interstellar silicate emission near 10\\mum~ in the direction of the Trapezium. This was confirmed by Becklin \\etal (\\cite{Becklin}) who also noticed extended silicate emission around $\\theta^2$~Ori~A. Since that time, interstellar silicate emission has been found by the Infrared Space Observatory (ISO) in the \\HII~ region N 66 of the Small Magellanic Cloud (Contursi \\etal in preparation) and in a few Galactic compact \\HII~ regions (Cox \\etal in preparation) and Photodissociation Regions (PDRs, Jones \\etal in preparation). The emission consists of two broad bands centered at 9.7 and 18\\mum, which show little structure and are clearly dominated by amorphous silicates.  We report in the present article ISO observations of the Orion bar and of a part of the Orion nebula made with the Circular Variable Filter of the ISO camera (CAM-CVF) which allowed imaging spectrophotometry of a field $3\\arcmin\\times3\\arcmin$ at low wavelength resolution (R $\\simeq$ 40). We also use an ISO Short--Wavelength Spectrometer (SWS) observation which provides higher--resolution ($R\\simeq 1000$) spectroscopy at a position within the \\HII~ region (see Fig. 2). This spectrum was taken as part of the MPEWARM guaranteed-time program. We show here that these new ISO data confirm and extend previous observations of the amorphous silicate emission and also give evidence for emission by crystalline silicates.  Sect. 2 of this paper describes the observations and data reduction. In Sect. 3, we discuss the emission of dust and gas. The silicate emission is characterized in section 4 through modelling of the observed continuum IR emission. Finally, conclusions are presented in Sect. 5. Our observations also give information on the fine--structure lines and on the AIBs. This will be presented in Appendices A and B respectively. ",
        "conclusions": "We obtained a rather complete view of the infrared emission of the Orion nebula and its interface with the adjacent molecular cloud. The most interesting results are the observation of amorphous, and possibly crystalline, silicates in emission over the entire \\HII~ region and in an extended region around the bright O9.5Vpe star $\\theta^2$ Ori A. We have fitted the mid--IR continuum of the \\HII~ region and around $\\theta^2$ Ori A with the emission from amorphous silicate and amorphous carbon grains at the equilibrium temperatures predicted for the grains in the given radiation field.  This shows that both types of grains can survive in the harsh conditions of the \\HII~ region. A number of bands (the 9.6\\mum~bump seen in Fig. 6; the excess 14\\mum~emission indicated in Figs. 4 and 10) suggest emission from crystalline silicates (essentially forsterite) in the \\HII~ region.  Crystalline silicates may also exist around $\\theta^2$ Ori A, but further, longer wavelength observations are required to confirm their presence. \\\\  Do the observed crystalline silicates result from processing of amorphous silicates in the \\HII~ region or in the environment of $\\theta^2$ Ori A?  Silicate annealing into a crystalline form requires temperatures of the order of 1000~K for extended periods (Hallenbeck \\etal \\cite{Hallenbeck}). The dust temperatures observed in the \\HII~ region and around $\\theta^2$ Ori A are considerably lower than this annealing temperature. One might however invoke grain heating following grain--grain collisions in the shock waves that are likely to be present in the \\HII~ region. However, grain fragmentation rather than melting is the more likely outcome of such collisions (Jones \\etal \\cite{JTH}). It is probable that the crystalline silicates observed here were already present in the parent molecular cloud, and probably originate from oxygen--rich red giants. \\\\  Emission by both amorphous and crystalline silicates has been observed with ISO around evolved stars (Waters \\etal \\cite{Waters}; Voors \\etal \\cite{Voors}). The crystalline silicates there must have been produced locally by annealing of amorphous silicates. Gail \\& Sedlmayr (\\cite{Gail}) have shown that this is possible, and that both amorphous and crystalline forms can be released into the interstellar medium. However, there is no evidence for absorption by crystalline silicates in the general interstellar medium in front of the deeply embedded objects for which amorphous silicate absorption is very strong (Demyk \\etal 1999, Dartois \\etal \\cite{Dartois}). Consequently, crystalline silicates represent only a minor fraction compared to amorphous silicates. It would be difficult to detect the emission from a small crystalline component of dust in the diffuse interstellar medium because the dust is too cool (T$\\,\\sim 20\\,$K) to emit strongly in the $15-40\\,\\mu$m wavelength region. Observations of \\HII~ regions and bright stars provide the opportunity of observing this emission due to the strong heating of dust.  Emission from amorphous and crystalline silicates is seen around young stars (Waelkens \\etal \\cite{Waelkens}; Malfait \\etal \\cite{Malfait}) as well as in comets (Crovisier \\etal \\cite{Crovisier}). There are also silicates in meteorites, but their origin is difficult to determine because of secondary processing in the solar system. Crystalline silicates in comets, and perhaps in interplanetary dust particles believed to come from comets (Bradley \\etal \\cite{Bradley}), must be interstellar since the material  in comets never reached high temperatures. However, the silicates probably experienced changes during their time in the interstellar medium. It is interesting to note that while very small grains of carbonaceous material exist, there seem to be no very small silicate grains in the interstellar medium (D\\'esert \\etal 1986)."
    },
    "0002/hep-ph0002076_arXiv.txt": {
        "abstract": "{ At the end of inflation, the coherent oscillations of the inflaton field may resonantly amplify the long wavelength modes of both bosons and fermions coupled to it. We study the resonant production of both kinds of particles during preheating in a model of hybrid inflation. The coherent time evolution of the inflaton and the Higgs fields after inflation induce a very different production of fermions depending on whether they are coupled to the Higgs or to the inflaton. For reasonable values of the model parameters, the fermion production through parametric resonance can be very efficient. We study the relative growth of the fermion and boson energy densities during preheating in hybrid models. During the initial stage of preheating, fermion production dominates the relative energy density, while the exponential growth of bosonic modes soon takes over.} ",
        "introduction": "Cosmological inflation is an extremely efficient mechanism in diluting any parti\\-cle species or fluctuations. At the end of inflation, the universe is empty and extremely cold, dominated by the homogeneous coherent mode of the inflaton. Its potential energy density is converted into particles, as the inflaton field oscillates  around the minimum of its potential. These particles are initially very far from equilibrium, but they interact among themselves (or decay) and thermal equilibrium is achieved at a very large temperature. From there on the universe expanded isoentropically, cooling down as it expanded, in the way described by the standard hot Big Bang model. Recent developments in the theory of reheating suggest that the decay of the inflaton energy could have been explosive~\\cite{KLS}, due to the coherent oscillations of the inflaton zero mode, which induce its stimulated decay. The result is a resonant production of particles in just a few inflaton oscillations, in a process known as preheating~\\cite{KLS}. The number of particles produced in this way is exponentially large, which may account for the extraordinarily large entropy, of order $10^{89}$, in our observable patch of the universe today.  Preheating is not generic, it may occur in different models of inflation~\\cite{GKLS,GBL,GKS} but only under special circumstances. Preheating strongly depends on the inflaton couplings to other fields, as well as on the amplitude and frequency of the inflaton oscillations. This explosive production occurs through parametric resonance of the long wavelength modes of {\\em any} field coupled to the inflaton, either bosonic~\\cite{KLS} or fermionic~\\cite{GK}. Due to the parametric nature of the resonance the particle production occurs for well defined frequency bands. The difference between bosonic or fermionic preheating is spectral: bosons can have large occupation numbers for a single mode, while fermions saturate due of the Pauli exclusion principle, and therefore a larger fraction of their energy can be transferred to higher momentum modes. Also, the rate of growth of these two components is very different: while the boson energy density grows exponentially~\\cite{KLS}, that of fermions cannot~\\cite{BHP}. As a consequence, very soon bosons dominate the decay of the inflaton energy if both kinds of particles are resonantly produced.  After a few inflaton oscillations, the energy density in fermions and bosons may be large enough to backreact on the inflaton oscillations and hence this will eventually stop their resonant production~\\cite{GKLS,GRTP}.  After preheating in chaotic inflation, the main part of the energy density of the universe is in the coherent modes of bosons: the inflaton, as well as any other parametrically amplified bosonic fields with large occupation numbers, while a subdominant part may be in the parametrically produced fermions. This state is very far from thermal equilibrium (characterized by a decoherent ensemble with small occupation numbers for all momentum modes) and, in fact, such a state could be used to generate the required baryon asymmetry during preheating, either via GUT baryogenesis~\\cite{GUTB}, EW baryogenesis~\\cite{EWB} or leptogenesis~\\cite{GRTP}. At the end, the final thermalization of the universe occurs through the rescattering of all the particles, which breaks the coherence of the bosonic fields. See, for instance, Ref.~\\cite{AS} for thermalization via fermionic modes.  In this paper we study the resonant production of both fermions and bosons in a hybrid model of inflation~\\cite{hybrid}. In these models it is possible to choose, without fine-tuning, the masses and couplings of the fields in such a way that the rate of expansion is negligible compared to the masses involved~\\cite{GBW}. In that case, preheating occurs in a few inflaton oscillations after the end of inflation, before the scale factor has grown by a single e-fold, and therefore we can ignore the universe expansion. This reduces the problem to resonant production of fermions in Minkowski space, for which there is a complete analytic treatment~\\cite{BHP}. We find that in the case of hybrid inflation, contrary to what happens in chaotic inflation, it is possible to produce, for certain couplings of the fermions to the inflaton field, a larger contribution of fermionic modes than bosonic ones to the energy density of the universe at the end of preheating.  We describe in Section 2 the hybrid model under consideration, with a Higgs-type field coupled to the inflaton, and discuss the solution of the classical equations of motion.  In Section 3 we study the pair production of fermions in this model, for a non-zero coupling of fermions to both the inflaton and the Higgs field. We analyze the growth of the fermion energy density for different values of the couplings, as well as for different fermion masses. In Section 4 we study the boson production for both couplings to inflaton and Higgs, together with the growth of the bosonic energy density.  In Section 5 we draw our conclusions. ",
        "conclusions": "In this paper we have studied the parametric resonant production of fermions in hybrid inflation, with both fields, inflaton and Higgs coherently oscillating after inflation. We have assumed that fermions may couple to either the inflaton or the Higgs (or both). The behaviour in the two cases is very different. While fermion production is very important in the case of a coupling to the inflaton, even in the presence of a bare mass, the production of fermions coupled only to the Higgs is generically weak. This is related to the fact that the non-adiabaticity condition, $d\\Omega_k/d\\tau >\\Omega_k^2$, is harder to achieve for the Higgs since, when $\\Omega_k$ is at a minimum, $d\\Omega_k/d\\tau$ is also at a minimum (contrary to the inflaton case). When the bare fermion mass exceeds the value $\\bar m_\\psi >\\sqrt{q_1}$ the fermion production by the coupling to the inflaton field is also suppressed.  We have studied the growth of the fermion energy density and seen that it very quickly saturates to an approximately constant value. For fermions coupled to the inflaton, the asymptotic value grows with the resonance parameter like $h_1^2 q_1^{1/4}$. For natural values of the couplings, a significant fraction of the inflaton energy can be transferred to fermions.  On the other hand, for fermions coupled to the Higgs, the fermion energy density is of the order $\\rho_{_F}/\\rho_{\\rm total} \\sim 10^{-2}g^4/\\lambda$, which under our working assumptions ($g^2\\ll\\lambda$) is quite small.  We have also studied the boson production for both couplings to inflaton and Higgs, and compared with the fermion production. While the energy density transferred to the parametrically produced fermions saturates after a few oscillations, the one in bosons grows exponentially with time. Hence, if both fermions and bosons have similar couplings, most of the particle production goes initially into fermions, while at late times the boson production is exponentially dominant. There are however some values of boson couplings to the Higgs, corresponding to $q'_2=n(n+1)/2$,  which completely inhibit the parametric resonance of bosons.  When the energy density of the bosonic or fermionic particles produced becomes sizeable, they are expected to backreact on the inflaton, affecting its evolution and eventually suppressing the parametric production of particles. We have not considered this process in detail in this paper since the approach followed is not the most suited one for this purpose. For a proper discussion of this issue in the context of chaotic inflation, see Refs.~\\cite{GRTP,BHP}.   As a summary, we have shown that the production of fermions in the preheating stage of hybrid inflation can be very important. For the range of model parameters assumed, hybrid inflation models lead to a more efficient fermion production than chaotic inflation models \\cite{GRTP}, without the need to go to extremely large values of the resonance parameter $q$.  Depending on the relative size of the couplings, and on the backreaction process, the inflaton energy transferred to fermions may even be larger than that transferred to bosons. At any rate, the parametric production of out of equilibrium fermions could have interesting consequences for cosmological issues such as the generation of the baryon asymmetry through e.g. the leptogenesis mechanism.  Note added: while completing the writing of this work a related paper \\cite{BGM} appeared where the gravitino production during preheating is computed in a supersymmetric model of hybrid inflation."
    },
    "0002/astro-ph0002296_arXiv.txt": {
        "abstract": "The physics of the standard hot big bang cosmology ensures that the early Universe was a primordial nuclear reactor, synthesizing the light nuclides (D, \\3he, \\4he, and \\7li) in the first 20 minutes of its evolution.  After an overview of nucleosynthesis in the standard model (SBBN), the primordial abundance yields will be presented, followed by a status report (intended to stimulate further discussion during this symposium) on the progress along the road from observational data to inferred primordial abundances.  Theory will be confronted with observations to assess the consistency of SBBN and to constrain cosmology and particle physics.  Some of the issues/problems key to SBBN in the new millenium will be highlighted, along with a wish list to challenge theorists and observers alike. ",
        "introduction": "Among the quantitative, ``hard\" sciences, astronomy has traditionally been scorned, with particular disdain reserved for cosmology.  No more. In the decade of the nineties the combination of an avalanche of high quality observational data and theoretical advances driven by enhaced computer (and brain) power, have succeeded in transforming cosmology to a precise science.  In this introductory lecture to IAU Symposium 198 on The Light Elements and Their Evolution it is my intent to describe primordial nucleosynthesis in this precision era of cosmology and to highlight the challenges, along with some goals, for the new millennium. After a brief review of the important physics during the era of primordial nucleosynthesis in the standard, hot big bang cosmological model (SBBN), I will present an overview of the predicted primordial abundances, emphasizing the generally very small theoretical uncertainties.  These will then be compared to the present best estimates (including their uncertainties) of the primordial abundances inferred from current observational data.  After assessing the consistency of SBBN, I will explore what SBBN has to offer to Cosmology and to Particle Physics and, what Cosmology may teach us about SBBN.  I will conclude with a summary of the key issues/problems confronting SBBN and with a wish list of topics I hope will be addressed during this meeting -- and beyond. ",
        "conclusions": "The study of the early evolution of the Universe and, in particular primordial nucleosynthesis, has truly entered the precision era of cosmology.  Precise abundances of the light nuclides are predicted and inferred from observations and the two are -- apparently -- in excellent agreement. As pleased as we may be at this success, it behooves us to avoid the temptation to rest on our laurels and to test this consistency ever more carefully.  To this end, it doesn't take much contemplation to identify several clouds looming on the horizon.  What follows is my personal list of some problems/issues I would like to see addressed at this Symposium and beyond.  \\subsection{Problems/Issues}  First consider deuterium.  On the one hand, any determination of the D/H ratio anywhere, anytime provides a {\\it lower} bound to the primordial abundance.  On the other hand, since ``wrong\" velocity hydrogen can masquerade as deuterium, any observation of ``deuterium\" is really an {\\it upper} bound to its true abundance.  More data tracking the velocity structure of the absorbing features used to identify D and H and exploring variations in D/H in material with similar histories will be very valuable.  More data at high-redshift and low-metallicity will be very valuable. After all, at present we are drawing profound conclusions on the basis of only two such systems.  Much remains to be done concerning the primordial abundance of \\4he.  For the most part, the \\hii regions from which the helium abundance is inferred have been modelled as homogeneous spheres or plane-parallel slabs.  A glance at the beautiful HST images of real \\hii regions reveals that they are anything but such idealizations.  What are the effects of temperature and/or density inhomogeneities, and how large may they be?  What of underlying stellar absorption which, if present but neglected, would lead to an {\\it under}estimate of the helium abundance.  And, what of the usually neglected ionization correction for neutral hydrogen and helium (Viegas, Gruenwald \\& Steigman 1999; see Viegas \\& Gruenwald, this volume)?  Considering this latter work, where models of \\hii regions ionized by realistic spectra of young star clusters were used in a reanalysis of the IT data, a {\\it reduction} in \\Yp of order 0.003 was suggested.  A comparison with Figure 1 shows that if \\Yp were reduced by this amount, the overlap between theory and data would, in fact, disappear.  \\subsection{Wish List}  Given the setting of this Symposium (Natal) and the proximity to the Christmas season, I'd like to conclude with my personal wish list.  To avoid being greedy, I'll only ask for two gifts.  A half-dozen or so observations of deuterium in high-z, low-Z systems along the lines-of-sight to distant quasars, with D/H determined in each (on average) to 10\\% or better.  With such a gift, I could determine $\\eta$ to better than 4\\%, predict \\Yp to $\\la 0.0007$, and constrain $\\Delta N_{\\nu}$ to an uncertainty less than $\\pm$~0.05.  I'd be a very happy cosmologist indeed.  My second wish is for \\4he measured to 3\\% accuracy (or better) in each of about a dozen low-metallicity, extragalactic \\hii regions, with care taken to address the several problems outlined above.  With such data, \\Yp could be fixed to better than the current level of $\\pm 0.002$, permitting \\4he to be used as a baryometer ($\\Delta\\eta/\\eta~\\la 20\\%$)."
    },
    "0002/astro-ph0002069_arXiv.txt": {
        "abstract": "Taking optically thin accretion flows as an example, we investigate the dynamics and the emergent spectra of accretion flows with different outer boundary conditions (OBCs) and find that OBC plays an important role in accretion disk model. This is because the accretion equations describing the behavior of accretion flows are a set of {\\em differential} equations, therefore, accretion is intrinsically an initial-value problem. We argue that optically thick accretion flow should also show OBC-dependent behavior. The result means that we should seriously consider the initial physical state of the accretion flow such as its angular momentum and its temperature. An application example to Sgr A$^*$ is presented. ",
        "introduction": "It has long been assuming that the parameters describing the accretion flow include the accretion rate, the mass of the central black hole, the viscosity parameter, and the parameter describing the strength of the magnetic field in the accretion flow. Once these parameters are given, we can obtain almost all the information of the accretion flow including the dynamics and the emergent spectrum. However, the set of equations describing the accretion flow are nonlinear differential equations, therefore it is intrinsically an initial-value problem. Thus the outer boundary condition (OBC) possibly plays an important role.  On the other hand, the complicated astrophysical environments make the physical states of the accreting gas at the outer boundary $r_{\\rm out}$, such as its temperature and angular momentum, various. For example, in semi-detached binary system, where the critical Roche lobe is filled up and the accretion of matter takes place through the inner Lagrangian point, the angular momentum of the accreted gas should be high; while in detached binary system the accretion matter is stellar winds therefore their angular momenta are much lower (Illarionov \\& Sunyaev 1975). In the nuclei of galaxies, where the supply of the accretion matter is  unclear, the initial physical states of the accretion flows should be more complicated. The complexity of astrophysical environments makes it important to investigate the role of OBC in accretion disk model. ",
        "conclusions": ""
    },
    "0002/astro-ph0002319_arXiv.txt": {
        "abstract": "We present He-like line ratios (resonance, intercombination and forbidden lines) for totally and partially photoionized media. For solar plasmas, these line ratios are already widely used for density and temperature diagnostics of coronal (collisional) plasmas. In the case of  totally and partially photoionized plasmas, He-like line ratios allow for the determination of the ionization processes involved in the plasma (photoionization with or without an additional collisional ionization process), as well as the density and the electronic temperature. \\\\ With the new generation of X-ray satellites, Chandra / AXAF, XMM and Astro-E, it will be feasible to obtain both high spectral resolution and high sensitivity observations. Thus in the coming years, the ratios of these three components will be measurable for a large number of non-solar objects.\\\\ In particular, these ratios could be applied to the Warm Absorber-Emitter, commonly present in Active Galactic Nuclei (AGN). A better understanding of the Warm Absorber connection to other regions (Broad Line Region, Narrow Line Region) in AGN (Seyferts type-1 and type-2, low- and high-redshift quasars...) will be an important key to obtaining strong constraints on unified schemes.\\\\ We have calculated He-like line ratios, for $Z$=6, 7, 8, 10, 12 and 14, taking into account the upper level radiative cascades which we have computed for radiative and dielectronic recombinations and collisional excitation. The atomic data are tabulated over a wide range of temperatures in order to be used for interpreting a large variety of astrophysical plasmas.\\\\ ",
        "introduction": "\\label{sec:Introduction} The new X-ray satellites (Chandra, XMM and Astro-E) will offer unprecedented high spectral resolution and high sensitivity spectra. Indeed, it will be possible to observe and to separate, in the X-ray range, the three most intense lines of He-like ions: the {\\bf resonance} line ({\\bf w}: 1s$^{2}$\\,\\element[][1]{S}$_{\\mathrm{0}}$ -- 1s2p\\,\\element[][1]{P}$_{\\mathrm{1}}$), the {\\bf intercombination} lines ({\\bf x,y}: 1s$^{2}$\\,\\element[][1]{S}$_{0}$ -- 1s2p \\element[][3]{P}$_{2,1}$ respectively) and the {\\bf forbidden} line ({\\bf z}: 1s$^{2}$\\,\\element[][1]{S}$_{0}$ -- 1s2s\\,\\element[][3]{S}$_{1}$). They correspond to transitions between the $n$=2 shell and the $n$=1 ground state shell (see Figure~\\ref{gotrian}).\\\\ The ratios of these lines are already widely used for collisional (coronal) plasma diagnostics of various types of objects: solar flares, supernovae remnants, the interstellar medium and tokamak plasmas, i.e. for very hot collisional plasmas (Mewe \\& Schrijver \\cite{Mewe78a} \\cite{Mewe78b}, Winkler et al. \\cite{Winkler81}, Doyle \\& Schwob \\cite{Doyle82}, and Pradhan \\& Shull \\cite{Pradhan81}). As shown by Gabriel \\& Jordan (\\cite{Gabriel69}, \\cite{GabrielJordan72}, \\cite{Gabriel73}), these ratios are sensitive to electron density (R(n$_{\\mathrm{e}}$), equation~\\ref{eq:R}) and to electronic temperature (G(T$_{\\mathrm{e}}$), equation~\\ref{eq:G}): \\begin{equation} \\mathrm{ R~(n_e)~=~\\frac{z}{(x+y)}} \\label{eq:R} \\end{equation} \\begin{equation} \\mathrm{ G~(T_e)=\\frac{z+(x+y)}{w}} \\label{eq:G} \\end{equation} As emphasized by Pradhan (\\cite{Pradhan85}), Liedahl (\\cite{Liedahl99}) and Mewe (\\cite{Mewe99}) (see also Paerels et al. \\cite{Paerelsetal98}), these plasma diagnostics could be also extended to study photoionized plasmas. Indeed, Pradhan has calculated the {\\bf R} and {\\bf G} ratios for highly charged ions (\\ion{Ar}{xvii} and \\ion{Fe}{xxv}) in ``recombination dominated non-coronal plasmas''. We present numerical calculations of these ratios, for six lighter ions, which could be applied directly for the first time to Chandra and XMM observations of the Warm Absorber present in Active Galactic Nuclei (AGN), and especially in Seyfert\\,1.\\\\  \\indent The Warm Absorber (WA) is a totally or a partially photoionized medium (with or without an additional ionization process), first proposed by Halpern (\\cite{Halpern84}) in order to explain the shape of the X-ray spectrum of the \\object{QSO MR2251-178}, observed with the Einstein Observatory. Its main signatures are the two high-ionization oxygen absorption edges, \\ion{O}{vii} and \\ion{O}{viii} at 0.74 keV and 0.87 keV respectively, seen in fifty percent of Seyfert 1 galaxies at least (Nandra $\\&$ Pounds \\cite{Nandra94}, Reynolds \\cite{Reynolds97}, George et al. \\cite{George98}). According to Netzer (\\cite{Netzer93}), an emission line spectrum from the WA should also be observed. Indeed, He-like ion lines have been observed in different types of Seyfert galaxies (\\object{NGC 3783}: George et al. \\cite{George95}, \\object{MCG-6-30-15}: Otani et al. \\cite{Otani96}, \\object{E 1615+061}: Piro et al. \\cite{Piro97}, \\object{NGC 4151}: Leighly et al. {\\cite{Leighly97}, \\object{NGC 1068}: Ueno et al. {\\cite{Ueno94}, Netzer \\& Turner \\cite{Netzer97}, and Iwasawa et al. \\cite{Iwasawa97}). The WA is supposed to be at least a two-zone medium with an inner part (called the ``inner WA'') associated with \\ion{O}{viii} and an outer part (called the ``outer WA''), less ionized, associated with \\ion{O}{vii} (Reynolds \\cite{Reynolds97}, Porquet et al. \\cite{Porquet99}). Furthermore, the \\ion{O}{vii} line is predicted to be the strongest line associated with the outer WA; the \\ion{Ne}{ix} line is predicted to be one of the strongest lines formed in the inner WA (Porquet et al. \\cite{Porquet98}).\\\\ The ionization processes, that occur in the Warm Absorber, are still not very well known. Indeed, even though the WA is commonly thought to be a photoionized gas, an additional ionization process cannot be ruled out (Porquet \\& Dumont \\cite{PorquetDumont98}, Porquet et al. \\cite{Porquet99}, Nicastro et al. \\cite{Nicastro99}). Thus, in the present paper, we do not restrict ourselves to only a single type of plasma, but rather study the following cases.\\\\ % We consider a ``pure photoionized plasma'' to be a plasma ionized by high energy photons (external ionizing source). For such a plasma, H-like radiative recombination (and dielectronic recombination at high temperature) are dominant compared to electronic excitation from the ground level (1s$^{2}$) of He-like ions. The lines are formed by recombination.\\\\ A ``hybrid plasma'' is a partially photoionized plasma, but with an additional ionization process, e.g. collisional (internal ionizing source). For this case, He-like electronic excitation processes from the ground level are usually as important as H-like recombinations, and may even dominate. The lines are formed by collisional excitation from the ground level with or without recombination.\\\\  In the next section, we introduce the atomic data calculations needed for such plasmas and we emphasize the role of upper-level radiative cascade contributions calculated in this paper for the populations of the $n$=2 shell levels. In section~\\ref{sec:diagnostics}, we develop line diagnostics of the ionization process (temperature) and the density for pure photoionized and hybrid plasmas. We give the corresponding numerical calculations of the line ratios for \\ion{C}{v}, \\ion{N}{vi}, \\ion{O}{vii}, \\ion{Ne}{ix}, \\ion{Mg}{xi}, and \\ion{Si}{xiii}. In section~\\ref{sec:practical}, we give a practical method for using these results to determine the physical parameters of the WA, in the context of the expected data from the new X-ray satellites (section~\\ref{sec:satellite}). ",
        "conclusions": "\\label{sec:conclusion} We have shown that the ratios of the three main lines (forbidden, intercombination and resonance) of He-like ions provide very powerful diagnostics for totally or partially photoionized media. For the first time, these diagnostics can be applied to non solar plasmas thanks to the high spectral resolution and the high sensitivity of the new X-ray satellites Chandra/AXAF, XMM and Astro-E.\\\\ These diagnostics have strong advantages. The lines are emitted by the same ionization stage of one element, thus eliminating any uncertainties due to elemental abundances. In addition, since the line energies are relatively close together, this minimizes wavelength dependent instrumental calibration uncertainties, thus ensuring that observed photon count rates can be used almost directly.\\\\ For example, the determination of the physical parameters of the Warm Absorber component in AGN, such as the ionization process, the density and in some case the electronic temperature (``hybrid plasma''), will allow observers to deduce the size and the location (from the ionizing source) of the WA. In addition, since He-like ions are sensitive to different range of parameters (density, temperature), it could permit confirmation of the idea that the WA comes from a stratified, or a multi-zone medium (Reynolds \\cite{Reynolds97}, Porquet et al. \\cite{Porquet99}). As a consequence, a better understanding of the WA will be important for relating the WA to other regions (Broad Line Region, Narrow Line Region) in different AGN classes (Seyferts type-1 and type-2, low- and high-redshift quasars...). This will offer strong constraints on unified schemes."
    },
    "0002/astro-ph0002405_arXiv.txt": {
        "abstract": "We apply the reconstruction technique of Nusser \\& Dekel (1992) to the recently available Point Source Catalogue Redshift Survey (PSCz) in order to subtract the phase correlations that are expected to develop in the mild non-linear regime of gravitational evolution. We study the evolution of isodensity contours defined using an adaptive smoothing algorithm in order to minimize the problems derived from the non-comutivity of operators. We study the genus curves of these isodensity contours and concentrate on the evolution of the amplitude drops, a meta-statistic able to quatify the level of phase-correlation present in the field. In order to test the method and to quatify the level of statistical uncertainty, we apply the method to a set of mock PSCz catalogues derived from the N-body simulations of two 'standard' CDM models, kindly granted to us by the Virgo consortium. We find the method to be reliable in recovering the right amplitude drops. When applied to PSCz the level of phase correlations observed is very low on all scales ranging from $5\\lu$ to $60\\lu$, providing support to the theory that structure originated from gaussian initial conditions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002155_arXiv.txt": {
        "abstract": "We here investigate the possibility that the ultra-high energy  cosmic ray (UHECR) events observed above the GZK limit are mostly protons accelerated in reconnection sites just above the magnetosphere of newborn millisecond pulsars which are originated by accretion induced collapse (AIC). We formulate the requirements for the acceleration mechanism and show that AIC-pulsars with surface magnetic fields $10^{12} $ G  $<  \\, B_{\\star} \\lesssim  $ $10^{15}$ G and spin periods 1 ms $\\lesssim \\, P_{\\star} \\, < \\, $ 60  ms, are able to accelerate particles to energies $\\geq \\, 10^{20} $ eV. Because the expected rate of AIC sources in our Galaxy is very small ($\\sim \\,  10^{-5}$ yr$^{-1}$), the corresponding contribution to the flux of UHECRs is  neglegible, and the total flux is given  by the integrated contribution from  AIC sources produced by the distribution of galaxies located within the distance which is unaffected by the GZK cutoff ($\\sim \\, 50 $ Mpc). We find that reconnection should convert a fraction $ \\xi \\, \\gtrsim \\, 0.1$ of magnetic energy into UHECR in order to reproduce the observed flux. ",
        "introduction": "The detection of  cosmic ray events with energies beyond 10$^{20}$eV by AGASA (Takeda et al. 1999), Fly's Eye (Bird et al. 1995), and Haverah Park (Lawrence, Reid, \\& Watson 1991) experiments still poses a challenge for the understanding of their nature and sources. These ultra-high energy cosmic rays (UHECRs) show no major differences in their air shower characteristics to cosmic rays at lower energies and thus one would expect them to be mostly protons (Protheroe 1999). If UHECRs are charged particles, or protons, then they should be affected by the expected Greisen-Zatsepin-Kuzmin (GZK) energy cutoff ($\\sim 5\\times 10^{19}$ eV), which is due to photomeson production by interactions with the cosmic microwave background radiation, unless they are originated at distances closer than about 50 Mpc  (e.g., Protheroe \\& Johnson 1995, Medina Tanco, de Gouveia Dal Pino \\& Horvath 1997). On the other hand, if the UHECRs are mostly protons from nearby sources (located within $\\sim $ 50 Mpc), then the arrival directions of the events should point toward their sources since they are expected to be little deflected by the intergalactic and Galactic magnetic fields (e.g., Stanev 1997, Medina Tanco, de Gouveia Dal Pino \\& Horvath 1998). The present data shows no significant large-scale anisotropy in the distribution related to the Galactic disk or the local distribution of galaxies, although some clusters of events seem to point to the supergalactic plane (Takeda et al. 1999, Medina Tanco 1998).  A number of source candidates and acceleration mechanisms have been invoked but all of them have their shortcomings (see, e.g., Olinto 2000 for a review). Particles can, in principle, extract the required energies from an induced e.m.f. in a circuit connected between the polar and the last open field line of a rapidly rotating pulsar, although it is not clear how the large voltages  can be maintained (e.g., Hillas 1984), or be accelerated in reconnection sites of magnetic loops, $if$ these can be produced, e.g., by Parker instability, on the surface of a pulsar (Medina Tanco, de Gouveia Dal Pino \\& Horvath 1997), but the accelerated particles will probably lose most of their energy gain by curvature radiation while dragged along by the magnetic dipole  field (Sorrell 1987).  Alternatively, Olinto, Epstein, \\& Blasi (1999) have recently proposed that UHECRs could be iron nuclei stripped by strong electric fields from the surface of highly magnetized neutron stars and accelerated in a relativistic MHD wind.  It is not clear however,  how the accelerated particles can escape from the magnetosphere of the star, for although efficient power extraction may be possible, there is a dense positron-electron plasma generated with the relativistic wind (Gallant \\& Arons 1994) that will possibly modify the electric fields and also interact with the energetic particles.  Their model also predicts that a correlation with the Galactic plane should become evident as data collection at the highest energies improves. In this paper we discuss an alternative model in which UHECRs are accelerated in magnetic reconnection sites outside the magnetosphere of  very young millisecond pulsars being produced by accretion induced collapse (AIC) of a white dwarf. ",
        "conclusions": "We have discussed the possibility that the UHECR events observed above the GZK limit are protons accelerated in reconnection sites just above the magnetosphere of very young  millisecond pulsars originated by accretion induced collapse. AIC-pulsars with surface magnetic fields $10^{12} $ G  $<  \\, B_{\\star} \\lesssim  $ $10^{15}$ G and spin periods 1 ms $\\lesssim \\, P_{\\star} \\, < \\, $ 60  ms, are able to accelerate particles to energies $\\geq \\, 10^{20} $ eV. These limits can be summarized by the condition $B_{\\star} \\,  \\gtrsim \\, 10^{13}$ G $  (P_{\\star}/2.5 $ ms)$^{4/3}$ (Eq. 1 and Fig. 2). Because the expected rate of AIC sources in our Galaxy is very small, the total flux is given  by the integrated contribution from  AIC sources produced in the distribution of galaxies within a volume which is unaffected by  the GZK cutoff (of radius $R_G \\simeq 50 $ Mpc). We find that the reconnection efficiency factor needs to be $ \\xi \\, \\gtrsim \\, 0.1$ in order to reproduce the observed flux. This result is appealing because it predicts no correlation of UHECR events with the Galactic plane, in agreement with present observations. However,  as data collection improves, we should expect some sign of correlation with the local distribution of galaxies and the supergalactic plane. These predictions can be ultimately tested by coming experiments such as the AUGER Observatory which will provide high statistics samples of UHECRs.    The model predicts a highly super-Eddington accretion mass rate  during part of the AIC process. In such a regime, one should expect a large mass outflow from the stellar surface itself which might alter the coronal conditions near the helmet streamer region. However, the strong closed magnetic fields in the pulsar magnetosphere can inhibit the gas outflow from the stellar surface if the surface temperature does not exceed the value at which the radiation energy density  ($a \\, T^4$) equals the magnetic field energy density ($B_{\\star}^2/8 \\pi$), i.e., $T_{\\star} \\lesssim 5 \\times 10^9 $ K $ B_{13}^{1/2}$; this upper limit is consistent with the predicted values for a neutron star that is formed by AIC (Woosley \\& Baron 1992).  Let us consider now the energy loss mechanisms that may affect the efficiency of the acceleration of the particles to the UHEs. Energy losses by curvature radiation which  occur when cosmic rays have to stream along magnetic field lines with a finite radius of curvature $R_c$, have a  cooling time $t_{rc} \\propto  R_c^2 $. In our model, the reconnection takes place in the helmet streamer region (see Fig. 1) where  $R_c \\rightarrow \\infty$ and therefore, $t_{rc} > > t_a$, where $t_a$ is the acceleration time for a proton in the reconnection site. [We can use the dimensions  of the reconnection site derived in \\S 2 to estimate the order of magnitude of $t_a \\simeq \\Delta R_X/ c \\, \\simeq 8 \\times 10^{-6} {\\rm s} \\, \\left({\\frac{\\Delta R_X/R_X}{0.1} }\\right) \\, \\left({ \\frac{R_X} {2.45 \\times 10^6  {\\rm cm} } }\\right)$.] The time for synchrotron losses  is $t_{syn} \\propto \\theta^{-2}$, where $\\theta$ is the particle pitch angle. Since we require that the protons in the reconnection zone are accelerated by an induced electric field rather than by a scattering process, they are expected to maintain beamlike pitch angles while escaping along the magnetic field lines, i.e., $\\theta < < 1$  and $t_{syn} > > t_a$.  The accelerated protons may also undergo  energy loss by pion and $e^{\\pm}$ pair production due to interactions with photons from the accretion disk radiation field. The characteristic distance scales for these processes can be estimated by $\\lambda_{p \\gamma} \\simeq (\\sigma_{p \\gamma} n_{\\gamma})^{-1}$  and $\\lambda_{pair} \\simeq (\\sigma_{pair} n_{\\gamma})^{-1}$ for photo-pion  and pair production, respectively, where $n_{\\gamma}$ is the photon number density at the reconnection site, $\\sigma_{p \\gamma} \\simeq  2.5 \\times 10^{-28} $ cm$^{-2}$ is the cross section for pion production, and $\\sigma_{pair} \\simeq  10^{-26} $ cm$^{-2}$ is the cross section for pair production (e.g., Bednarek and Protheroe 1999). In order to estimate $n_{\\gamma}$, we must determine the luminosity in the accretion disk. At super-Eddington accretion rates, the disk is thicker and heat is radially advected with matter.  The luminosity of such advective supercritical accreting disks is given by $L_D \\simeq [0.6 + 0.7 \\, ln (\\dot M_D/\\dot M_{Edd})] \\, L_{Edd}$ (e.g., Lipunova 1999), where $L_{Edd} \\simeq 1.25 \\times 10^{38} $ erg s$^{-1}$ $(M_{\\star}/ M_{\\odot})$. Using the value obtained in \\S 2 for $\\dot M_D/ \\dot M_{Edd} \\simeq 1.8 \\times 10^{9}$, we find $L_D \\simeq 1.8 \\times 10^{39}$ erg s$^{-1}$. Inside the disk, where the optical depth is much larger than unity, photons and particles are in thermodynamic equilibrium and the disk radiates like a black-body with a temperature $T_D \\simeq 1.1 \\times 10^9 $ K $\\left({ \\frac {\\dot M_D } {3 \\times 10^{-8} M_{\\odot} \\, {\\rm s}^{- 1} } }\\right)^{1/4} \\, \\left({ \\frac{ M_{\\star}} {M_{\\odot} } }\\right)^{1/4} \\, \\left({ \\frac{ R_X} {2.45 \\times 10^6  {\\rm cm} } }\\right)^{-3/4}$. At a distance $R_X$, $n_{\\gamma} \\simeq L_D/ (\\pi R_X^2 \\bar \\epsilon_{\\gamma})$ (e.g., Sorrell 1987), where $\\bar \\epsilon_{\\gamma} \\, \\simeq \\, 2.8 k_B T_D \\sim 0.35 MeV$ is the mean photon energy, or $n_{\\gamma} \\simeq 8 \\times 10^{21} {\\rm cm}^{-3} \\, \\left({ \\frac{ L_D} { 1.9 \\times 10^{39} {\\rm erg} \\, {\\rm s}^{-1} } }\\right) \\, \\left({ \\frac{ R_X} {2.45 \\times 10^{6}  {\\rm cm} } }\\right)^{-2} \\, \\left({ \\frac{ T_D} { 1.5 \\times 10^9  {\\rm K} } }\\right)^{-1}$. We note that beyond the disk, the luminosity can be substantially  smaller as the matter just above the disk surface  can be partially opaque to the outgoing radiation. Thus, the value above of $n_{\\gamma} $ for the helmet streamer is probably overestimated. Substitution of $n_{\\gamma}$ into the equation for $\\lambda_{p \\gamma}$ gives $\\lambda_{p \\gamma}/R_X  \\, \\gtrsim  \\, 0.22 \\, > 0.1 \\left({\\frac{\\Delta R_X/R_X } { 0.1}}\\right)$, so that the interaction distance scale for pion production is larger than the width of the acceleration region. For pair production, although the interaction distance $\\lambda_{pair} $ is shorter, the inelasticity of this process is so small (inelasticity $K_p \\sim \\, 2 \\times 10^{-3}$) that the corresponding energy loss rates of the UHECR protons are even smaller than in the case of pion production (e.g., Bednarek \\& Protheroe 1999). However, the injected $e^{\\pm}$ pairs can initiate inverse Compton pair cascades in the reconnection region that may saturate an induced electric field if the width of the reconnection region is larger than $\\lambda_{sat} \\simeq 100 \\, (\\sigma_{trip} n_{\\gamma})^{-1}$, where $\\sigma_{trip} \\simeq  1.3 \\times 10^{-26} $ cm$^{-2}$ is the cross section for the triple $e^{\\pm}$ pair production by electron-photon collisions. We find that $\\lambda_{sat}/R_X \\gtrsim 0.4 > 0.1 \\left({\\frac{\\Delta R_X/R_X } { 0.1}}\\right)$, so that the width of the reconnection site is smaller than $\\lambda_{sat}$."
    },
    "0002/astro-ph0002225_arXiv.txt": {
        "abstract": "We review the evidence that detectable fluxes of X-rays are produced by inverse Compton scattering of nuclear photons with the relativistic electrons in the radio lobes of strong FRII radio galaxies within the FRII-RL quasar unification scheme. We report here on the possible detection of this effect in two steep spectrum RL quasars. This may have important implications on the physics and evolution of powerful radio galaxies. ",
        "introduction": "It is well known that the synchrotron radio emission from the extended lobes of strong radio galaxies and radio-loud (RL) quasars samples ultra relativistic electrons. It is a customary practice to estimate the average magnetic field intensity via the minimum energy argument by making use of the emitted radio flux in the 10 MHz -- 100 GHz band (source rest frame). Since the critical frequency emitted in the synchrotron process is $\\nu(MHz) \\sim B(\\mu G) (\\gamma/10^3)^2$ and typically $B(\\mu G) > 10$, it then follows that only electrons with Lorentz factor $\\gamma > 10^3$ are taken into account. These electrons are also responsible for producing X-rays via the inverse Compton (IC) scattering of the cosmic microwave background (CMB) photons [$\\epsilon(keV) \\sim (\\gamma/10^3)^2$]. When these X-rays are detected, the number of ultra-relativistic electrons is fixed and from the radio flux one can uniquely determine the average magnetic field intensity. Up to now, because of the weakness and diffuse nature of the predicted X-ray emission, detection of X-ray fluxes due to this process has been possible for a few sources only, notably Fornax A (Feigelson et al. 1995; Kaneda et al. 1995) and Cen B (Tashiro et al. 1998); the derived magnetic field intensities are lower than the classical equipartition values by factors 1.5--2.  Brunetti, Setti \\& Comastri (1997) have pointed out that sizeable X-ray fluxes can also be emitted by the IC scattering of relativistic electrons in the FRII's radio lobes with the IR photons from a quasar, and associated circumnuclear dusty/molecular torus, hidden in the galaxy's nucleus. Since the IR emission peaks at $50-100 \\mu m$ electrons at lower energies ($\\gamma < 500$) are involved in this process. Of course there are no reasons why these lower energy particles shouldn't be present, on the contrary one would expect them on physical grounds based on acceleration and loss mechanisms. In order to estimate the size of the expected X-ray flux, for a given quasar IR emission, one may work out the equipartition by extrapolating downward the electron spectrum derived from the synchrotron emission to a minimum energy ($\\gamma_{min}$) limited from below by possible Coulomb losses. The equipartition fields ($B_{eq}$) so derived are stronger than the classical one by factors from 1.5 to 3 (also Setti, Brunetti \\& Comastri 1999). We have shown that the IC scattering of the IR nuclear photons may easily account for a large fraction of the extended X-ray emission of several powerful FRIIs at large redshifts ($z \\sim 1$) detected by ROSAT in the 0.1--2.4 keV interval.  Morphologically there are two important aspects that should be mentioned: firstly, for obvious geometrical reasons, the X-rays from the IC scattering of the quasar photons are more concentrated toward the nuclear region than those from the IC scattering of the CMB photons and, secondly, given two symmetrical radio lobes the X-ray emission from the far lobe can be much larger than that from the near one, depending on the orientation of the radio axis with respect to the line of sight, due to the enhanced efficiency of head-on scatterings (Brunetti et al.1997). It should also be mentioned that, while the X-rays from the IC scattering of the CMB photons must have a spectral slope coincident with that of the synchrotron radio emission, the X-ray spectrum associated with the IC scattering of the nuclear photons may or may not have the synchrotron slope simply because a different portion of the primary electron spectrum is being sampled (see also Brunetti 2000).  Direct evidence of extended X-ray emission from the IC scattering of the IR photons from a hidden quasar has been gathered by Brunetti et al.(1999) making use of ROSAT HRI observations of the powerful, double lobed radio galaxy 3C 219. The residual X-ray distribution after subtraction of the absorbed, unresolved nuclear source is remarkably coincident with the radio structure. The central extended ($\\sim$ 100 kpc) X-ray emission, somewhat stronger in the counter-jet side as expected in our model, can be accounted for by assuming a magnetic field $\\sim 3$ times weaker than our equipartion value ($B_{eq} \\simeq 10\\mu G$, $\\gamma_{min} = 50$). Of course this estimate depends on the assumed IR power of the hidden quasar which we have derived by two, albeit indirect, approaches since 3C 219 has not been detected by IRAS and not observed by ISO. Observations with {\\it Chandra}, scheduled in the fall of the year 2000, will likely provide a check of our model.    \\begin{figure*} \\centerline{ \\psfig{figure=3c215.ps,width=11cm,angle=270} } \\end{figure*} ",
        "conclusions": "There is supportive observational evidence of extended X-ray emission from the IC scattering of quasar IR photons with relativistic electrons in the lobes of powerful radio galaxies. Besides being a confirmation of the FRII--quasar unification, this may provide an important tool for the diagnostic of the relativistic plasma at particle energies not sampled by radio observations. The unavoidable presence of lower energy particles implies stronger magnetic fields than derived by standard equipartition formulae and, consequently, a larger pressure inside the lobes. Moreover, accounting for the extended X-ray emission in sources for which the quasar radiation can be constrained indicates magnetic field strengths lower than equipartition. Therefore, confirmation of our IC model by {\\it Chandra} and XMM satellites may provide important clues on the physics and evolution of radio sources."
    },
    "0002/astro-ph0002013_arXiv.txt": {
        "abstract": "These lectures cover various aspects of the statistical description of cosmological density fields. Observationally, this consists of the point process defined by galaxies, and the challenge is to relate this to the continuous density field generated by gravitational instability in dark matter. The main topics discussed are (1) nonlinear structure in CDM models; (2) statistical measures of clustering; (3) redshift-space distortions; (4) small-scale clustering and bias. The overall message is optimistic, in that simple assumptions for where galaxies should form in the mass density field allow one to understand the systematic differences between galaxy data and the predictions of CDM models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002019_arXiv.txt": {
        "abstract": "We present a detailed investigation of atmospheres around accreting neutron stars with high magnetic field ($B\\gtrsim 10^{12}$ G) and low luminosity ($L\\lesssim 10^{33}$  erg/s). We compute the atmospheric structure, intensity and emergent spectrum for a plane--parallel, pure hydrogen medium by solving the transfer equations for the normal modes coupled to the hydrostatic and energy balance equations. The hard tail found in previous investigations for accreting, non--magnetic neutron stars with comparable luminosity is suppressed and the X--ray spectrum, although still harder than a blackbody at the star effective temperature, is nearly planckian in shape. Spectra from accreting atmospheres, both with high and low fields, are found to exhibit a significant excess at optical wavelengths above the Rayleigh--Jeans tail of the X--ray continuum. ",
        "introduction": "The problem of investigating the properties of radiation emitted by neutron stars (NSs) accreting at low rates, $\\dot M\\approx 10^{10}-10^{14}$ g/s, became of interest after it was realized that the Galaxy may contain a large population of low luminosity magnetic accretors (see e.g. Nelson, Wang, Salpeter, \\& Wasserman \\cite{nel95:1995}). The Galaxy should harbor more than $10^3$ Be/X--ray binaries with an accreting NS shining at $\\approx 10^{32}-10^{34}$ erg/s (Rappaport, \\& van den Heuvel \\cite{rvdh82:1982}; van den Heuvel, \\& Rappaport \\cite{vdhr86:1986}). Moreover, assuming a supernova birth rate of 10--100 yr$^{-1}$, $\\sim 10^8-10^9$ old, isolated NSs (ONSs) should be present in the Galaxy. Accretion onto a strongly magnetized, moving neutron star may be severely hindered for different reasons, but there is the possibility that a small, albeit non--negligible, fraction of ONSs may be accreting directly from the interstellar medium and some of them might be above the sensitivity threshold of ROSAT (e.g. Treves, \\& Colpi \\cite{tc91:1991}; Blaes, \\& Madau \\cite{bm93:1993}; see Treves, Turolla, Zane, \\& Colpi \\cite{review:1999} for a review).  At variance with neutron stars accreting at high rates, e.g. in X--ray pulsators, in low--luminosity sources the interaction of the escaping radiation with the inflowing material is of little importance, so they provide a much simpler case for investigating the physics of accretion in a strongly magnetized environment. For luminosities far below the Eddington limit, the accretion problem becomes germane to that of calculating the spectrum emerging from static atmospheres around cooling NSs. Spectra from cooling NSs have been widely investigated by a number of authors in connection with the X--ray emission from young, millisecond pulsars and isolated NSs, both for low and high magnetic fields and for different chemical compositions (see e.g. Romani \\cite{rom87:1987}; Shibanov, Zavlin, Pavlov \\& Ventura \\cite{sh92:1992}; Rajagopal, \\& Romani \\cite{rr96:1996}; Pavlov, Zavlin, Tr\\\"umper, \\& Neuh\\\"auser \\cite{pztn96:1996}). Emerging spectra are not very different from a blackbody at the star effective temperature, the distinctive hardening present at low fields ($B\\lesssim 10^9$ G) becoming less pronounced when the magnetic field is $\\sim 10^{12}-10^{13}$ G. Similar conclusions were reached for the spectrum emitted by low--luminosity, low--field, accreting NSs by Zampieri, Turolla, Zane, \\& Treves (\\cite{ztzt95:1995}, hereafter ZTZT) for a pure hydrogen atmospheric composition.  The search of isolated neutron stars with ROSAT produced in recent years half a dozen promising candidates (Walter, Wolk, \\& Ne\\\"uhauser \\cite{wwn96:1996}; Haberl \\etal \\ \\cite{ha97:1997}; Haberl, Motch \\& Pietsch \\cite{hmp98:1998}; Schwope \\etal \\ \\cite{sch99:1999}; Motch \\etal \\ \\cite{mo99:1999}; Haberl, Pietsch \\& Motch \\cite{hpm99:1999}). All of them show a soft, thermal X--ray spectrum, with typical energies $\\sim 100$ eV, and have an exceedingly large X--ray to optical flux ratio, $\\gtrsim 10^4$. Although their association with isolated neutron stars is firmly established, their interpretation in terms of an accreting or a cooling object is still a matter of lively debate. Present models predict rather similar spectral distributions in both cases, especially in the X--ray band. It is therefore of particular importance to improve our theoretical understanding of these two classes of sources, looking, in particular, for spectral signatures which can enable us to discriminate between them.  In this paper we present a first detailed calculation of spectra emitted by strongly magnetized ($B\\gtrsim 10^{12}$ G), accreting neutron stars, focusing our attention on low luminosities, $L \\sim 10^{30}$--$10^{33}$ erg/s, such as those expected from old neutrons stars accreting the interstellar medium. Spectral distributions are computed solving the transfer equations for the normal modes coupled to the hydrostatic equilibrium and the energy balance for different values of the accretion rate and the magnetic field.  We find that spectra emerging from magnetized, accretion atmospheres are blackbody--like in the X--ray band, in close agreement with the known results for cooling, magnetized atmospheres. However, accretion spectra show a new and distinctive feature at low energies, being characterized by an excess over the Raleigh--Jeans tail of the X--ray continuum below $\\sim 10$ eV. The same behaviour is found in accretion atmospheres around unmagnetized neutron stars, but, as already pointed out by ZTZT, the X--ray spectrum is sensibly harder in this case. This result may be relevant in connection with the isolated neutron star candidate RX J18563.5-3754. Multiwavelength observations of this source indicate that, while ROSAT data are well fitted by a blackbody at $T_{eff}\\sim 60$ eV, HST points lie above the extrapolation of the fit in the optical (Walter, \\& Matthews \\cite{wm97:1997}).   The plan of the paper is as follows. The input physics relevant to our model is presented in \\S \\ref{model}: radiative transfer in a magnetized plasma is discussed in \\S \\ref{model-radtra} and the structure of an accreting, magnetized atmosphere in \\S \\ref{model-struct}. Computed spectra are presented in \\S \\ref{results}. Discussion and conclusions follow in \\S \\ref{discuss}. ",
        "conclusions": "\\label{discuss}  We have discussed the spectral distribution of the radiation emitted by a static, plane--parallel atmosphere around a strongly magnetized neutron star which is heated by accretion. Synthetic spectra have been computed solving the full transfer problem in a magnetoactive plasma for several values of the accretion luminosity and of the star magnetic field. In particular, we explored the low luminosities ($L\\sim 10^{30}$--$10^{33}$ erg/s), typical e.g. of isolated accreting NSs, and found that model spectra show a distinctive excess at low energies over the blackbody distribution which best--fits the X--rays. The energy $\\hbar\\omega_{ex}$ at which the spectrum rises depends both on the field strength and the luminosity, but for $B\\approx 10^{12}$--$10^{13}$ G $\\hbar\\omega_{ex}\\sim 10$ eV, so that the optical emission of an accreting, NS is enhanced with respect to what is expected extrapolating the X--ray spectrum to optical wavelengths. The presence of an optical/UV bump is due to the fact that the low--frequency radiation decouples at very low values of the scattering depth in layers where the gas is kept at larger temperatures by Compton heating. Since a hot outer zone is exhibited by accreting and not by cooling atmospheres, the optical excess becomes a distinctive spectral feature of NSs accreting at low rates.  Despite present results are useful in shedding some light on the emission properties of accreting, magnetized neutron stars, many points still need further clarification before the problem is fully understood and some of the assumptions on which our investigation was based deserve a further discussion. In this paper we have considered only fully ionized, pure hydrogen atmospheres. In the case of cooling NSs, the emitted spectra are strongly influenced by the chemical composition of the surface layers which results from the supernova explosion and the subsequent envelope fallback. Present uncertainties motivated several authors to compute cooling spectra for different abundances (see e.g. Miller, \\& Neuh\\\"auser \\cite{mn91:1991}; Miller \\cite{mil92:1992}; Pavlov, \\etal \\ \\cite{pszm95:1995}; Rajagopal, Romani, \\& Miller \\cite{rrm97:1997}) and led to the suggestion that the comparison between observed and synthetic X--ray spectra may probe the chemistry of the NS crust (Pavlov, \\etal \\ \\cite{pztn96:1996}). The assumption of a pure hydrogen atmosphere, although crude, is not unreasonable for an accretion atmosphere. In this case, in fact, incoming protons and spallation by energetic particles in the magnetosphere may enrich the NS surface with light elements (mainly H) and, owing to the rapid sedimentation, these elements should dominate the photospheric layers (Bildsten, Salpeter, \\& Wasserman \\cite{bsw92:1992}).  Even in this simplified picture, however, the assumption of complete ionization may not provide an entirely realistic description. In a strongly magnetized plasma atomic binding is greatly enhanced, mainly for light elements. For a typical field of $\\sim 10^{12}$ G, the ionization potentials for the hydrogen ground state are $\\sim 100-300$ eV, moving the photoionization thresholds into the soft X--rays (see e.g. Ruderman \\cite{r72:1972}; Shibanov \\etal \\ \\cite{sh92:1992}). The role of photoionization and of pressure ionization in a magnetized hydrogen atmosphere has been investigated by a number of authors (e.g Ventura, Herold, Ruder, \\& Geyer \\cite{vhrg92:1992}; Potekhin, \\& Pavlov \\cite{pp97:1997}; Potekhin, Shibanov, \\& Ventura \\cite{psv98:1998}). The neutral fraction, $f_H$, reaches a peak at densities $\\approx 1 \\ {\\rm g/cm}^3$ then decreases due to pressure ionization and turns out to be highly dependent on the temperature. While $f_H$ never exceeds a few percent at $T = 10^6$ K, it becomes as large as $\\sim 80 \\%$ for $T = 10^{5.5}$ K and $B=10^{13}$ G. Although in the models we presented here such low temperatures are only reached in relatively low--density layers (see figure \\ref{tempfig}), hydrogen ionization equilibrium should be properly included in a more detailed analysis.  The effects of different orientations of the magnetic field through the atmosphere have been also neglected by keeping $\\bf{B}\\, ||\\, \\bf{n}$, a key simplifying assumption which allowed us to solve the transfer problem in one spatial dimension. Clearly, such an approximate description is valid only if the size of the emitting caps is small. Due to the intrinsic anisotropy of a magnetized medium, the emerging flux is expected to depend on the angle $\\theta_B$ between the magnetic field and the normal to the surface. Shibanov \\etal \\ (\\cite{sh92:1992}) estimated that in a cooling atmosphere the flux at 1 keV from a surface element perpendicular to the field may exceed that of an element parallel to ${\\bf B}$ by nearly 50\\%. This result shows that a non--uniform magnetic field may lead to a significant distortion in the emerging spectra. Moreover, if the orientation of $B$ varies along the atmosphere, tangential components of the radiative flux may induce some meridional circulation to maintain the heat balance (Kaminker, Pavlov, \\& Shibanov \\cite{kps82:1982}).  When the finite size of the emitting regions is accounted for, the problem can not reduced to a plane--parallel geometry and its solution necessarily demands for multidimensional transfer algorithms (e.g. Burnard, Klein, \\& Arons \\cite{bka88:1988}, \\cite{bka90:1990}; Hsu, Arons, \\& Klein \\cite{hak97:1997}). Calculations performed so far were aimed to the solution of the frequency--dependent radiative problem on a fixed background or of the full radiation hydrodynamical problem in the frequency--integrated case. Numerical codes for the solution of the full transfer problem in axially symmetric media under general conditions are now available, see e.g. ZEUS (Stone, \\& Norman \\cite{sn92:1992}), ALTAIR (Dykema, Castor, \\& Klein \\cite{dck96:1996}) and RADICAL (Dullemond, \\& Turolla \\cite{dt99:1999}). Their application to the transfer of radiation in accretion atmospheres around magnetized NSs can add new insights on the properties of the emitted spectra."
    },
    "0002/astro-ph0002475_arXiv.txt": {
        "abstract": "ESA and NASA are studying projects having a tremendous return on variable star research. Other national space agencies are also studying or developing projects of smaller costs but with impressive returns. The projects range from global Galactic surveys like the ESA mission GAIA which will give photometric time series for about 1 billion stars, to detailed pulsation modes studies like the CNES mission COROT which could reach a photometric precision lower than 1 ppm. The presentation will emphasize the future astrometric, asteroseismologic and planet detection missions. ",
        "introduction": "Many advantages are present when going in space. The most exploited one has been the access to wavelengths not transmitted through the atmosphere. For photometry near the visible part of the spectrum, there are many benefiting aspects: No atmosphere means no absorption bands of $H_2O$ and $O_3$, no scintillation, no emission or diffusion from the sky (dark sky), no variable extinction. Furthermore depending on the satellite orbit, it may be possible to make long continuous monitorings (however eclipses by the Moon or the Earth may occur) or  to make a global sky coverage.  As a comparative example, the COROT team (see section~6) studied fictive Earth projects (cf. Baglin 1997a) and showed that even an immoderate project of 8 8m-class telescopes disseminated around the Earth would not surpass their small space mission.  The drawbacks of space projects are their complexity, cost, aging, limited operation time, and many technical constraints. Space missions are fragile on technical and also political point of views. The global budgets of ESA and NASA in local currencies have been relatively flat for about 8 years, but they may be subjected to high volatility according to changes in the policy. ",
        "conclusions": "Many of these projects might be selected in a near future for their realisation. The projects performing astrometry will provide deep galactic surveys and thus an inventory for galactic variable stars. The asteroseismological projects have common goals but they use different ingenious techniques and have different and diverse by-products.  The future spatial projects will open new domains in variable star studies, therefore precise predictions are delicate. Since some presented missions are already funded, there is the conviction that after the projects like Hipparcos or the microlensing surveys, the domain of variable stars will still be under a strong evolution."
    },
    "0002/astro-ph0002369_arXiv.txt": {
        "abstract": "The measurement of cosmological parameters is investigated in a representation of the least-action method that uses a redshift-space dataset to simultaneously constrain the real-space fields $\\delta$,$\\b v$. This method is robust in recovering the entire evolution of the matter density contrast and peculiar velocities of galaxies in real space from current galaxy redshift surveys. The main strength of the method is that it permits us to break the degeneracy of the parameters $b$ and $\\Omegam$ (customarily measured in the ratio $\\beta\\equiv \\Omegam^{0.6}/b$ from redshift-space distortions), and these are evaluated in the current context separately. The procedure provides a simple numerical means to extract as much information as possible from a given sample, in the simplest linear bias model, before resorting to cosmic complementarity to resolve the degeneracy in the measurement of $\\Omegam$. The same premise applies to more sophisticated choices of bias models. We construct a likelihood parameter $\\lambda(b,\\Omegam)$ to evaluate the relative likelihood of different values of $b$ and $\\Omegam$. The method is applied to the \\iras~redshift survey with a low-resolution Gaussian smoothing length of 1200 \\km within a spherical region $x_{\\rm max} \\sim 15,000$ \\km and the reconstructed velocity field is then compared with POTENT-reconstructed velocities from the Mark III radial-velocity dataset within a radius $\\sim 5000$ \\km, which have been suitably prepared to account for Malmquist bias and other systematic errors. The analysis yields a likelihood for the parameters that is overall consistent with $\\Omegam\\approx 0.3$ and $b\\approx 1.1$, thus lending support to a non-vanishing cosmological constant $\\Omega_{\\Lambda}\\approx 0.7$ in a flat universe. ",
        "introduction": "\\begin{figure*} \\centering \\begin{picture}(210,210) \\special{psfile='fig0.ps' angle=-90 voffset=290 hoffset=-160 vscale=60 hscale=67} \\end{picture} \\caption[]{Qualitative distribution of errors in the reconstruction of the matter fields from redshift surveys. The error $\\epsilon_0$ in the current redshift sample increases monotonically in (a) as the perturbative solutions propagate $\\epsilon_0$ to increasing amplitudes when integrated back in time; in the boundary condition problem of the LAP method, shown in (b), errors fluctuate between the fixed end-points.} \\end{figure*}  Galaxy redshift surveys are undoubtedly extremely valuable tools to investigate the evolution of the universe at large scales. The cosmologist's prerogative is to determine the evolution of the matter density contrast $\\delta$ and peculiar velocity $\\b v$ that yields such cosmic structure, customarily assuming that it formed solely by gravity, and the cosmological parameters that determine their dynamics. In the standard paradigm of a FRW expanding universe, the interplay of both fields is governed by the density parameter $\\Omega_0$ and the Hubble parameter $H_0$. On the other hand, a relationship between the fields $\\delta$,$\\b v$ and the survey data is established by adopting a bias model that purports the correlation between the $z$-space galaxy number-count and the underlying matter field. Devoid of such a relationship, the edifice of measuring cosmological parameters from galaxy redshift surveys has no foundation whatsoever. A standard working hypothesis, that I shall accept throughout this paper, is that of linear bias, i.e. $b^2 \\equiv P(k)_{\\rm gals}/P(k)_{\\rm matter}$ (more elaborate bias models are propounded in e.g. Dekel \\& Lahav 1999). For simplicity we shall also leave out the scale-dependence of $b$. Therefore, three relevant parameters that are interesting to pin down from redshift-space samples are in this context $\\Omegam$,$H_0$ and $b$. In this paper I shall be chiefly concerned with $\\Omegam$ and $b$ ($H_0$ will be scaled out with distance).  Tracing back in time the matter fields takes us to an initial epoch of fluctuations of very small amplitude $\\delta\\lsim 10^{-4}$, seeded by a period of inflationary expansion. At that point the information derived from the galaxy surveys connects with early-universe data such as the spectrum of fluctuations on the CMB. If the matter fields could realistically be traced back to such a primordial stage by integrating the equations of gravitational instability, then the statistics of the $\\delta$ field would be a potentially key discriminant to rule out cosmological models. For instance, non-gaussianity in the initial $\\delta$ field rules out most inflationary models, and only those leading to a non-Gaussian primordial spectrum remain acceptable (such models are suggested in e.g. Linde, Sasaki \\& Tanaka 1999).  Kaiser (1987) proposed measuring cosmological parameters from redshift-space distortions by virtue of the fact that overdense regions appear to be flatter along the line-of-sight in redshift space. This distortion, quantified by the parameter $\\beta=\\Omegam^{0.6}/b$, permits us to solve the equations for $\\delta$,$\\b v$, at least perturbatively (see e.g. Dekel 1994; Coles \\& Sahni 1995), and measurements of $\\beta$ have been investigated in much detail in the literature (Strauss \\& Willick 1995; Dekel 1994,1999a; Dekel, Burstein \\& White 1997). Also, in view of the fact that the bias parameter is almost certainly dependent on the selected sample, estimates have been computed for $\\beta_{IRAS}$ given $b_I$ for {\\it IRAS} galaxies (Dekel \\etal 1993; Fisher \\etal 1995a; Willick \\etal 1997a,b; Sigad \\etal 1998; more recently from the PSC$z$ sample, Canavezes \\etal 1998; Tadros \\etal 1999; Saunders \\etal 2000) and from the Optical Redshift Survey (ORS) (Hudson \\etal 1995; Santiago \\etal 1995; Baker \\etal 1998). The Mark III peculiar velocity survey similarly yields estimates of $\\beta$ from redshift distortions (Willick \\etal 1995,1996,1997a,b; Dekel, Burstein \\& White 1997; Sigad \\etal 1998). It is only beyond the linear approximation (i.e. $\\delta \\propto \\nabla\\cdot\\b v$) and, indeed, beyond the assumption of linear bias, that one can break down the degeneracy between $\\Omegam$ and $b$ and estimate these parameters separately, rather than via $\\beta$ (Fry 1994; Bernardeau \\etal 1995). Verde \\etal (1998) achieved this by proposing the bispectrum as a measure of cosmological parameters, in a model of non-linear bias. In this paper we also pursue breaking the degeneracy of $\\Omegam$ and $b$ from the redshift-space data and show that by using the least-action framework it is indeed possible to do so within the linear bias model.  The least-action principle (LAP) was first used in the Local Group by Peebles (1989,1990). The trajectories of nearby galaxies were computed subject to two boundary conditions: vanishing initial velocities and fixed present positions. This simple scenario of self-gravitating point-like masses with two boundary conditions produced an estimate of $\\Omegam$ by fitting to the observations the predicted peculiar velocities of nearby galaxies. The LAP method has also been used as a test of $\\Omegam=1$ CDM models (Branchini \\& Carlberg 1994), as well as to integrate the orbits of a significant number of galaxies from partial coverage redshift samples (e.g. Shaya, Peebles \\& Tully 1995). An equivalent representation of the LAP method in terms of continuous fields, i.e. the density contrast and velocity fields was proposed by Giavalisco \\etal (1993), and employed in Susperregi \\& Binney (1994)(hereafter SB94) and Susperregi (1995) in the reconstruction of $\\Omegam=1$ simple models, such as exact solutions and Gaussian random fields. More recently, Schmoldt \\& Saha (1998) proposed a variant of the customary LAP formulation by rewriting the equations motion in redshift space.  The key difference between the variational and perturbative approaches lies on how the errors are spread over the time-reversed evolution. This is qualitatively sketched in Fig.~1. A $n$th-order solution differs, in the time-reversed direction, from the true solution by a monotonically growing parameter $\\epsilon$ which sets out from a small value $\\epsilon(t_0)$ (at any rate $\\epsilon_0$ is at least the sum of the systematic and random errors of the dataset) and the conservation of kinematical quantities is preserved up to $O(\\epsilon^{n})$. This is adequate within a time span $t_c \\ll t \\lsim t_0$ where $\\epsilon(t_c)\\sim 1$, and $t_c$ marks a transition into the loss of convergence. The distribution of errors in the LAP method on the other hand, is by construction evenly distributed along the trajectory; the initial and final boundary conditions are fixed, though not without systematic and numerical errors, and the parameter $\\epsilon$ fluctuates along the trajectory between both end-points (Fig.~1b). Hence the solution is well-behaved whether the errors remain within the bound $\\epsilon \\lsim 1$ or not. In that respect there is an advantage with respect to perturbative solutions; the downside of it is of course that within the span of time where perturbative solutions are valid, LAP errors may fluctuate with larger amplitude than the perturbative equivalent. The LAP method, in a nutshell, thus consists in finding Ansatze for the matter fields that optimize the distribution of $\\epsilon$ along the phase-space trajectory, and hence minimize the overall departure with respect to the exact solution. The following two difficulties may arise: \\begin{itemize} \\item {\\bf A} {\\it Finding ``dynamically plausible'' solutions}. If the matter field is sparsely sampled or the errors in the dataset are substantial, then the boundary condition given by the survey, taken at face value, may not correspond to the outcome of gravitational evolution from the initial fluctuations (typically $\\delta\\sim 0$ or vanishing peculiar velocities). The LAP method will in this case find a {\\it dynamically plausible} fit between the end-points, which will be as faithful a representation of the true evolution as is the quality of the dataset. \\item {\\bf B} {\\it Formation of multistreams in over-dense regions}. Multistreams are characterised by galaxies at the same redshift which are located at different positions along the line of sight and have different infalling velocities. The degeneracy in redshift among streams makes them indistinguishable and hence compatible but inequivalent solutions result, as many as there are streams. The LAP method cannot discriminate among these solutions; multistreams indeed erase the memory of their past evolution. \\end{itemize} The second problem can only be overcome by casting aside part of the information contained in the sample and smoothing over the existing non-linearities to transform the multivalued field into a single-valued one, typically with a smoothing length $\\sim 500 - 1000$ \\km. The resulting smoothed field is clearly a less resolved representation of the underlying galaxy orbits, albeit the only tractable one.  The advent of large galaxy redshift surveys strengthens the motivation to use the LAP method. Near all-sky redshift surveys, e.g. PSC$z$, \\iras~and ORS provide an excellent sky coverage (within a galactic latitude $|b|\\gsim 8^{\\circ}$ for {\\it IRAS} galaxies and $|b|\\gsim 20^{\\circ}$ for ORS), that may be extended further to cover the Zone of Avoidance via a Wiener reconstruction (Fisher \\etal 1995b; Zaroubi \\etal 1999). They are therefore a fairly thorough representation of the underlying matter density field. Obviously the greater number of galaxies in the sample the more accurate is the representation of the field, and this is best achieved with a redshift survey. Real-space datasets require Tully-Fisher distance calibrations of individual galaxies, and consequently the end result is a sparser sampling than is achieved with the same computational effort by measuring redshifts and angular coordinates. The goal of this paper is to exploit galaxy redshift surveys to the best effect and extract as much information from them as is possible; the main thesis put forward is the LAP method demonstrably breaks down the degeneracy in the determination of $\\Omegam$ and $b$. This entails very tangible advantages. On the one hand, the freedom to investigate those two parameters separately permits us not to take the idea of bias seriously. A form of bias will certainly always be present in one form or another so that we can make sense of the galaxy number-count with respect to the underlying matter field. However, whether that is a linear or non-linear bias, the more one dissociates this phenomenological relationship from our measurements of $\\Omegam$, the more credible those measurements will be. This is indeed what LAP does. On the other hand, the LAP method produces a reconstruction on the basis of the redshift-space sample alone, free of any proviso regarding the shape of the power spectrum. Assuming a given shape for $P(k)$ unduly overconstrains the system, as will the addition of other datasets.  In this article, I shall mainly apply the LAP method to the \\iras~survey and study the predicted values of $b$ and $\\Omegam$. The reconstructed \\iras~velocity field is then compared with the Mark III velocity sample to seek a fine-tuning of the parameters. A more thorough undertaking, in terms of the quality of the sample, is to apply the LAP method to PSC$z$, which is by a factor of 3 a more densely sampled survey than \\iras, and it will be interesting to tackle this in future work. The article is structured as follows: Section 2 describes the LAP method in some detail and how to find solutions that are consistent with a redshift-space dataset; in Section 3 we test the method with several {\\it IRAS} mock catalogues obtained via $n$-body simulations; in Section 4 we apply the method to the \\iras~galaxy redshift survey, optimizing the predicted velocities with the Mark III dataset; finally, in Section 5 we summarize the main conclusions. ",
        "conclusions": "The LAP method provides a practical means to break the degeneracy between $\\Omegam$ and $b$ in galaxy redshift surveys. The method is employed in the manner of a nonlinear constraint on the redshift-space dataset and, although in formulation it comes across as algebraically cumbersome, it is of considerable simplicity and efficiency from the numerical point of view. The method is sound in that it does not require an a priori approximation of the map $\\b x\\to\\b s$ to pin down the solution and it provides considerable freedom to ascribe relative importance to the data available, i.e. the initial and final endpoints, to which we wish to invariably assign greater weight than intermedate stages of which little or no data are available.  The method can prove significant to measure $\\Omegam$ in the latest largest samples, and extract the most accurate information prior to comparison with other datasets, such as the CMB radiation power spectrum and SN data. One important challenge for the future is to attain a better grasp of the concept of bias and this will be probably achieved via microlensing data and $n$-body simulations of the formation of galaxies and clusters from primordial fluctuations, rather than from galaxy redshift surveys. Once a model of bias is adopted on a sound footing, then clearly the LAP model is impeccable in producing an estimate of $\\Omegam$. In the simple linear bias model we have employed we have totally relegated any consideration of scale-dependence in $b$. This is a point I have deliberately omitted for simplicity. Thus, the estimates computed in this paper ought to be regarded qualitatively as weighted averages of the ``real'' $b$ over different scales, if indeed scale-dependent bias models are to be believed.  In this paper, we have employed the likelihood function (\\ref{likelihood}) to investigate the values of $b$,$\\Omegam$. Clearly this is not a unique choice. However, our choice is guided by the argument of relative convergence of the solutions, which is justifiably a reasonable criterion to get close to the ``real'' solutions. In view of the performance of the $\\lambda$ function in the reconstruction of the mock samples, this choice does not appear to be totally off the mark. A potential reason for concern could be the offset observed between the maxima of the likelihood functions and the real values of the parameters in the $n$-body simulations. However the recurrence of this offset in a predictable manner lends strength to the argument that it arises as a numerical fault that is easy to account for systematically in the analysis of the datasets. The reconstructions of the fields are, on the other hand, of considerable accuracy and no numerical defficiency or hindrance is observed. The application of the method to \\iras~predicts the parameters to be fairly accurately located in the immediate neighbourhood of the maxima $\\Omegam\\approx 0.3$ and $b\\approx 1.1$, which is found to be most compatible with the estimate of $\\beta$ given by Willick \\etal (1997a). In a flat universe such predicted values are perfectly consistent with a non-vanishing cosmological constant or a quintessence scalar field component. The likelihood examined in this way is only very slightly modified when the velocities predicted via the LAP method are finely-tuned with data from the Mark III sample. The shift of the predicted values is towards slightly greater values of the parameters but it remains comfortably consistent with the results obtained from \\iras."
    },
    "0002/astro-ph0002125_arXiv.txt": {
        "abstract": "Some QSOs show an abrupt, strong rise in polarization at  rest wavelength $\\sim750$~\\AA.  If this arises in the atmosphere of an accretion disk around a supermassive black hole, it may have diagnostic value.  In PG 1222+228, the polarization rise occurs at the wavelength of a sharp drop in flux. We examine and reject interpretations of this feature involving a high velocity outflow.  The observations agree with a model involving several intervening Lyman limit systems, two of which happen to coincide with the Lyman continuum polarization rise.  After correction for the Lyman limit absorption, the continuum shortward of $\\lambda 912$ is consistent with a typical power-law slope, $\\alpha \\approx -1.8.$  This violates the apparent pattern for the Lyman limit polarization rises to occur only in ``candidate Lyman edge QSOs''.  The corrected, polarized flux rises strongly at the wavelength of the polarization rise, resembling the case of PG 1630+377.  The rise in polarized flux places especially stringent requirements on models. ",
        "introduction": "Spectropolarimetric observations with the {\\em Hubble Space Telescope} (\\hst) have revealed an unexpected rise in linear polarization in several QSOs (see review by Koratkar \\& Blaes 1999).  These are radio quiet ``candidate Lyman edge QSOs'', in which the continuum flux drops rather rapidly at rest wavelengths $\\lambda < 1000$~\\AA\\ (Antonucci, Kinney, \\& Ford 1989; Koratkar, Kinney \\& Bohlin 1992).  Models of accretion disk atmospheres predicted a reduced polarization in the Lyman continuum because of a diminished contribution of electron scattering to the opacity (Laor, Netzer, \\& Piran 1990).  With this motivation, Impey \\etal\\ (1995) and Koratkar \\etal\\ (1995) used the Faint Object Spectrograph (FOS) on \\hst\\  to obtain ultraviolet spectropolarimetry of several QSOs with redshifts sufficient to bring the Lyman continuum within the observed wavelength band.  The surprising result, in several cases, was a rapid {\\em rise} in polarization in the Lyman continuum.  From values $\\lesssim 1$ percent in the optical and near ultraviolet, the observed polarization rises around rest wavelength 750~\\AA\\ to values $\\sim5$ percent in several objects, and to $\\sim20$ percent in PG 1630+377.  This phenomenon has inspired several attemps at explanation.  Blaes and Agol (1996) found that, for effective temperatures $\\teff \\approx 25,000$~K and low effective gravities, a polarization rise of up to $\\sim5$~percent at about the observed wavelength could occur naturally in QSO disk atmospheres.  This results from the interplay of electron scattering, bound-free opacity, and the temperature gradient in the atmosphere.  However, Shields, Wobus, and Husfeld (1998, hereinafter SWH) showed that  the effects of the relativistic transfer function destroy the agreement between this model and observation.  Beloborodov and Poutanen (1999) suggested a model involving Compton scattering in a corona or wind, but this model appears to have trouble giving the rapid rise in polarized flux observed in PG 1630+377 (Blaes \\& Shields 1999).  Lee and Blandford (1997) discussed the possible role of scattering by resonance lines of heavy elements (see Section 5).  SWH showed that, if the polarization is assumed to rise sharply at $\\lambda 912$ in the rest frame of the orbiting gas, then relativistic effects would naturally produce the wavelength dependence of the observed polarization.  This may offer a way of measuring the black hole spin, but the physical mechanism for the polarization rise remains unknown.  PG 1222+228 is a $\\rm B \\approx 15.5$ radio quiet QSO (Schmidt \\& Green 1983) whose polarization rise at $\\lambda \\approx 750$~\\AA\\ coincides with a sharp drop in flux (Fig. 1, 2).  Impey \\etal\\ (1995) noted this and attributed it to a coincidental Lyman limit system (LLS), corresponding to an identified absorption line system at z = 1.486.  However, the coincidence of a broad absorption feature with a polarization rise also is observed for broad absorption line (BAL) QSOs.  In these objects, outflowing gas at velocities $\\sim10^4~\\kmps$ produces blueshifted absorption troughs, typically seen in the resonance lines of H I, C IV, N V, O VI, Si IV, and sometimes Mg II (Weymann \\etal\\ 1991; Arav, Shlosman, \\& Weymann 1997). Spectropolarimetric observations (\\eg, Ogle 1997; Schmidt \\& Hines 1999; Ogle \\etal\\ 1999) often show a rise in polarization in the troughs, reaching values as high as $\\sim8$ to 10 percent from $\\sim1$ percent at unabsorbed wavelengths.  This is explained in terms of scattering of some of the continuum by an extended region that is not covered by the BAL flow (Hines \\& Wills 1995; Goodrich \\& Miller 1995; Cohen \\etal\\ 1995). This pattern resembles the polarization rise and flux drop in PG 1222+228.  This paper addresses two questions: (1)  Does the polarization rise in PG 1222+228 result from an intrinsic absorber, analagous to the situation in the BAL QSOs?  (2)  If the drop in flux in PG 1222+228 is an intervening LLS, what are the consequences of correcting the observed, polarized continuum for this absorption? ",
        "conclusions": "The Lyman continuum polarization rises are among the more puzzling recent observational discoveries concerning QSOs.  The wavelength dependence, rising rather abruptly at nearly the same rest wavelength in the several known cases, suggests a connection with the bound levels of atoms.  The proximity of the feature to $\\lambda 912$ further suggests an association with the Lyman edge of hydrogen.  The ad hoc model of SWH supports an association with the Lyman edge and raises the possiblity of confirming the presence of a relativistic disk and constraining its parameters. However, attempts to fit the feature with a physical model have encountered difficulties.  This is an important problem for QSO theory.  Are the reported polarization rises real?  The coincidence of the polarization rise in PG 1222+228 with a sharp drop in flux might raise the question of background problems with the FOS spectropolarimeter.  Impey \\etal\\ (1995) argue that the polarization is unlikely to be less than 2.7\\% around 2000 \\AA\\ under any reasonable assumption for the FOS background.  However, the degree of polarization in the far ultraviolet is uncertain by at least a factor two because of systematic errors involving background and scattered light in the FOS. The observed polarization rises in several QSOs occur at different observed wavelengths but similar rest wavelengths.  We are not aware of any polarization rises of this nature in FOS spectropolarimetry of BL Lac objects or stars.  There is an urgent need for a renewed capability for ultraviolet spectropolarimetry from space to confirm and extend the measurements.  Lyman continuum polarization rises have heretofore been associated with the candidate Lyman edge QSOs (see discussion by Koratkar \\etal\\ 1998).  Our results suggest that PG 1222+228 may not be a true member of this class. Measurements of the Lyman continuum polarization in additional QSOs are needed to clarify the frequency of occurence of the phenomenon and to look for correlations with features in the continuum flux, in the line intensities and polarization, and other properties.  Observations to shorter rest wavelengths are needed to determine whether the polarization falls or continues to rise.  Measurements of the time dependence of the polarization rises would be most interesting.  The emitting radius of an accretion disk would be light weeks.  Lee and Blandford (1995) considered a model for the far ultraviolet polarization rise of QSOs that did not involve the Lyman edge.  Noting that a number of resonance lines of heavy elements fall in the rest wavelength range where the polarization rises, they suggested that resonance scattering of the QSO continuum might produce the observed polarization.  Such a model could produce a rising polarized flux, since the polarization could be essentially zero at wavelengths without scattering contributions.  As noted above, the polarization rise in PG 1222+228 may be too steep for models involving an accretion disk.  In this case, alternative models such as resonance scattering may hold promise.  We note that the polarized flux spectrum of PG 1630+377 (Koratkar \\etal\\ 1995) shows a strong rise at the wavelength of the N V emission line.  The claim by Richards \\etal\\ (1999) that luminous QSOs have many intrinsic, narrow, high velocity C IV absorption systems has important implications. This would complicate the use of such systems to probe the evolution of galaxies and the intergalactic medium.  The mechanism for producing the absorbing clouds would add another challenge to the subject of outflows from QSOs.  Surveys of QSOs at different luminosities, to a uniform standard of signal-to-noise ratio, would allow confirmation of the claimed higher incidence of absorption in more luminous QSOs. Tests of the intrinsic nature of narrow absorptions have been summarized by Barlow, Hamann, \\& Sargent (1997).  These include time variability of the depth and profile of the lines and evidence for saturated but nonblack line profiles.  Detection of changes in the velocities of the lines would be most revealing.  Chemical abundances may be an indicator, given evidence for high abundances in the broad absorption and emission-line regions of QSOs (\\eg, Hamann \\& Ferland 1993).  If C IV systems in high and low luminosity QSOs have similar, mostly subsolar abundances, this might argue for their intervening nature."
    },
    "0002/astro-ph0002063_arXiv.txt": {
        "abstract": "NGC~4945 is one of the brightest Seyfert galaxies on the sky at 100 keV, but is completely absorbed below 10 keV, implying an optical depth of the absorber to electron scattering of a few;  its absorption column is probably the largest which still allows a direct view of the nucleus at hard X--ray energies.  Our observations of it with the Rossi X--ray Timing Explorer (RXTE) satellite confirm the large absorption, which for a simple phenomenological fit using an absorber with Solar abundances implies a column of $4.5^{+0.4}_{-0.4} \\times 10^{24}$ cm$^{-2}$. Using a a more realistic scenario (requiring Monte Carlo modeling of the scattering), we infer the optical depth to Thomson scattering of $\\sim 2.4$.  If such a scattering medium were to subtend a large solid angle from the nucleus, it should smear out any intrinsic hard X--ray variability on time scales shorter than the light travel time through it.  The rapid (with a time scale of $\\sim$ a day) hard X--ray variability of NGC~4945 we observed with the RXTE implies that the bulk of the extreme absorption in this object does {\\sl not} originate in a parsec-size, geometrically thick molecular torus.  Limits on the amount of scattered flux require that the optically thick material on parsec scales must be rather geometrically thin, subtending a half-angle $< 10^\\circ$.  This is only marginally consistent with the recent determinations of the obscuring column in hard X--rays, where only a quarter of Seyfert~2s have columns which are optically thick, and presents a problem in accounting for the Cosmic X--ray Background primarily with AGN possessing the geometry as that inferred by us.  The small solid angle of the obscuring material, together with the black hole mass (of $\\sim 1.4 \\times 10^6$ $M_{\\odot}$) from megamaser measurements, allows a robust determination of the source luminosity, which in turn implies that the source radiates at $\\sim 10$\\% of the Eddington limit. ",
        "introduction": "Our current best picture of nuclei of Seyfert galaxies includes the central source (i.e. black hole, accretion disk and broad line region) embedded within an optically thick molecular torus (cf. Antonucci \\& Miller 1985).  The object is classified as a Seyfert~1 for viewing directions which lie within the opening angle of the torus, so that there is a direct view of the nucleus, and as a Seyfert~2 for directions intersecting the obscuring material.  The torus absorbs the optical, UV and soft X--ray nuclear light, so the nucleus in Seyfert~2s can only be seen at these energies through scattered radiation.  One of the central questions still under debate in the context of the Unification Models of the two types of Seyferts is: how optically thick is this putative torus?  NGC~4945 is a nearby (3.7 Mpc;  Mauersberger et al. 1996) edge--on galaxy. It has strong starburst activity, producing intense IR emission concentrated in a compact nuclear region (Rice et al. 1988;  Brock et al. 1988), and a ``superwind'' outflowing along the minor axis of the galaxy (Heckman, Armus, \\& Miley 1990). It also has an active nucleus, first seen unambiguously in Ginga X--ray observations (Iwasawa et al. 1993), confirming the Seyfert~2-type classification.  These data showed a heavily obscured, strong hard X--ray source above 10 keV, confirmed by the CGRO OSSE observations (Done, Madejski, \\& Smith 1996) which in turn revealed that NGC~4945 is one of the brightest extragalactic sources in the sky at 100 keV! The absorbing column, a few $\\times$ $10^{24}$ cm$^{-2}$, is among the largest which still allows a direct view of the nucleus at hard X--ray energies. If such a scattering medium were to subtend a large solid angle from the nucleus, it would smear out any intrinsic hard X--ray variability on timescales shorter than the light travel time through it.  The presence of the Seyfert nucleus is further supported by the fact that the object is a megamaser source (detected in the H$_2$O bands) implying an edge--on geometry, but one of the key features which makes its study so important is that it is one of only four AGN where the black hole mass can be constrained (at $\\sim 1.4 \\times 10^6$ $M_{\\odot}$) from detailed mapping of the  megamaser spots (Greenhill, Moran, \\& Herrnstein 1997).  As such, it is one of a few unique sources where the luminosity in Eddington units can be reliably estimated.  Below we present the data from the RXTE, confirming the large absorbing column, but also revealing large amplitude hard X--ray variability on a time scale of days.  A distant absorber with an appreciable optical depth, subtending a large solid angle as seen by the nucleus, would smear out the rapid variability on time scales shorter than the light travel time through such an absorber.  Given our data, we conclude that the optically thick absorber cannot be both distant and geometrically thick. ",
        "conclusions": "The variability we see in NGC~4945 is then entirely compatible with that expected from the {\\sl intrinsic} source, with no significant scattered delay by a distant material.  If this absorber has an appreciable Compton thickness (as is the case here), and if it subtends a large solid angle to the X--ray source, then it should intercept and scatter a large fraction of the flux.  If it is also distant from the nucleus, the light travel time effects will ``wash out'' any intrinsic variability on time scales shorter than the light travel time through the absorber.  Conversely, a structure with much smaller scale height subtends a much smaller solid angle, making scattering less important. This also precludes the observed X--rays to be purely due to Compton reflection, as this would require a contrived geometry of the reflector with respect to the primary X--ray source:  the reflector would have to be located very close to a completely covered central source.  We note that by comparison, a well studied unabsorbed Seyfert~1 NGC~3516 showed X--ray variability with a similar fractional amplitude on $\\sim 10 \\times$ longer time scales (Nandra \\& Edelson 1999) than seen in NGC~4945.  While the black hole mass in NGC~3516 is not as well known, it is estimated to be $\\sim 10^7$ $M_{\\odot}$.  The fact that the ratio of variability time scales is roughly the same as the ratio of nuclear masses -- as expected for accreting black holes -- further supports our conclusion that the hard X--ray variability we see in NGC~4945 is intrinsic.  With the optical depth to electron scattering of a few, the shape of the absorption cutoff would be different than expected from pure photoelectric absorption, and the detailed shape of the emergent spectrum depends on the geometry.  We model this numerically with a Monte Carlo code as given in Krolik, Madau, \\& \\.Zycki (1994), assuming a torus with square cross-section where the half-angle subtended by the torus, $\\theta_0$, its optical depth to electron scattering $\\tau_{\\rm e}$, and the power law index of the incident energy spectrum $\\alpha$ are free parameters (the Comptonization cutoff is set to 100 keV as above).  The results of our fits (using the PCA data over the range of 3 -- 20 keV, and HEXTE data as above) are shown in Table 1, where the 90\\% confidence regions on $\\Gamma$ and $\\tau_{\\rm e}$ are typically 0.1.  A small scale height absorber ($\\theta_0 \\sim 10^\\circ$) gives $\\tau_{\\rm e} = 2.4$, compared to a large scale height ($\\theta_0 \\sim 80^\\circ$) which gives $\\tau_{\\rm e} = 2.1$. With an iron abundance of twice Solar, these fits change to $\\tau_{\\rm e} = 1.7$, and $\\tau_{\\rm e} = 1.5$, respectively.  (Since the photoelectric cutoff present in our data is mainly sensitive to the column density of Fe, larger-than-Solar abundance of Fe would make us overestimate the true absorbing column if we assumed Solar abundances, and vice-versa.) While statistically these might marginally favor the large scale height absorber, we consider that all the fits are probably equally likely given that modeling the spectrum with a fixed cutoff energy may introduce systematic uncertainties.  Our calculations include the Fe K emission line produced by the torus but we also include an additional Fe line (such as may be expected to arise in the photoionized scattering medium).  Those calculations also imply that large ($> 4 \\times$ Solar) Fe abundances can be excluded:  they would imply a stronger Fe K line than is seen in our data. In reality, this limit is probably more stringent, given the fact that at least a fraction of the Fe K line originates in a more extended region.  These Monte Carlo results also give the distribution of the number of scatterings which the photons undergo before reaching the observer positioned in the equatorial plane. This is key in determining the solid angle subtended by the optically thick absorber, and thus its vertical size scale.  Fig.~3 shows the fraction of the observed photons that underwent 0, 1, 2, 3, etc. scatterings before reaching the observer for eight values of $\\theta_0$ as discussed above (cf. Table~1), with the solid and dotted lines showing the results for Solar and twice Solar abundance of iron, respectively.  The fraction of photons which arrive without being scattered is 19\\% and 63\\% respectively for a ``thick'' ($\\theta_0 = 80^\\circ$) and ``skinny'' ($\\theta_0 = 10^\\circ$) torus.  For $2 \\times$ Solar abundance of Fe these numbers are 32\\% and 75\\%.  The data in Fig.~2 imply that fewer than 40\\% of the observed photons are scattered over path lengths longer than 1 light day, so the half-angle subtended by the optically thick material is less than $\\sim 10^\\circ$. This implies a rather small scale height, and perhaps is due to the same material which produces the H$_{2}$O maser emission (Greenhill et al. 1997).  The details of the geometry of Seyfert~2s are important in the assessment of their contribution to the Cosmic X--ray background, as the heavily absorbed AGN were postulated to make up the bulk of it (cf. Krolik et al. 1994; Madau, Ghisellini, \\& Fabian 1994; Comastri et al. 1995).  The value of $\\theta_0$ of $10^\\circ$ is marginally consistent with the torus geometry inferred from recent observations at hard X--ray energies.  These observations show that Seyfert~2s outnumber Seyfert~1s by a factor of 4:1, while more than a quarter of Seyfert~2s have a column which is optically thick (see Giommi et al. 1998, and Gilli, Risaliti, \\& Salvati 1999 and references therein). Assuming a single geometry for the Seyfert nuclei where the torus has a rectangular cross-section, we can attempt to reproduce these ratios by assuming that if the central flux barely ``grazes'' the torus, we classify the object as {\\sl any} Seyfert~2, while an object is an optically thick Seyfert~2 only if the line of sight encounters the entire radial distance in the torus.  In this context, requiring such 1:4:1 ratio would then imply that the torus is somewhat flattened, with outer radius of 6.5 $\\times$ its inner radius and equatorial height of 1.3 $\\times$ its inner radius.  In this scenario, all Seyfert~2s are then seen at angles smaller than $\\sim 50^\\circ$ from the plane, while the optically thick Seyfert~2s are confined to angles of $\\le 12^\\circ$.  We repeated the Monte--Carlo calculations with this rectangular geometry and find that the fraction of scattered photons is $\\sim$ 50\\%, still too large to match the observed hard X--ray variability. A further problem arises if this is indeed a universal geometry for all Seyferts as the Thomson depth of the absorber is large, $\\sim 1.7$ in the more restrictive $2 \\times$ Solar case.  This absorber reduces the unabsorbed flux $\\sim$ five-fold or more, and this is not consistent with the 1:4:1 ratio, by a large margin.  We thus conclude that a population of AGN with geometry very nearly that of NGC~4945 cannot make up the CXB. Instead, significantly larger fraction of the heavily obscured AGN is required (implying a large solid angle subtended by the absorber) than implied from the rather small $\\theta_0$ inferred by us.  One plausible scenario would have the local optically thick Seyfert~2s surrounded by absorbers that already collapsed to an accretion disk, while in the more distant objects -- in the earlier stages of evolution -- such absorbers had larger vertical extent.  Alternately, the absorption could come from a structure which is much closer to the nucleus.  The variability limits impose constraints on the amount of scattered flux that is lagged on time scales of more than 1 day, and thus they do not constrain the height of any structure which is $<< 1$ light day from the X--ray source (corresponding to a distance of $<< 10^4$ Schwarzschild radii for the mass of the black hole of $\\sim 10^6$ $M_{\\odot}$). While a very geometrically thick accretion disk cannot be ruled out from our data, there are considerable theoretical difficulties in maintaining a structure with a large height scale. It is far easier to envisage a structure with a small height scale such an accretion disk with outer regions somewhat thickened by instabilities resulting from radiation pressure warping (cf. Maloney, Begelman, \\& Pringle 1996).  With these arguments for the Thomson-thick absorber subtending a small solid angle in NGC~4945, we can now estimate the true luminosity of the source.  The Monte Carlo simulations show that the intrinsic flux must be substantially larger than that observed (by a factor of $e^{\\tau_{\\rm e}}$, or $\\times 11$ for Solar abundances) because any scattered nuclear photons are lost from the line of sight. The 1-500 keV flux, corrected for photoelectric absorption alone, is $\\sim 5 \\times 10^{-10}$ erg cm$^{-2}$ s$^{-1}$, so correcting for the Thomson opacity yields a 1-500 keV intrinsic X--ray luminosity of $10^{43}$ erg~s$^{-1}$.  Assuming that the thermal (opt/UV/EUV) emission from the accretion disk is roughly equal to the hard X--ray emission gives a total bolometric luminosity of the nucleus of $\\sim 2 \\times 10^{43}$ erg~s$^{-1}$.  With this, and $M_{\\rm BH}$ of $1.4 \\times 10^6$ $M_{\\odot}$, the source is radiating at $\\sim$ 10\\% of the Eddington luminosity;  even if the abundances are twice-Solar (which yields $\\tau_{\\rm e}$ of 1.7), $L/L_{\\rm Edd}$ is at least $\\sim$ 5\\%. (We note that similar $L/L_{\\rm Edd}$ was also inferred by Greenhill et al. (1997), but the discovery of the rapid hard X--ray variability allows us to determine the source luminosity more accurately, as now we know that relatively few photons are scattered {\\sl back} to the line of sight.)  NGC~4945 is one of the few AGN where this quantity can be calculated robustly, since the mass of the central object is {\\sl known}, although we are aware that the value of its central mass is not as accurate as for the famous megamaser NGC~4258;  the resulting uncertainty in the estimate of $L/L_{\\rm Edd}$ may be a factor of 2, comparable to the effects of the unknown Fe abundance or the ratio of $L_{\\rm Tot} /L_{\\rm X-ray}$.  The resulting $L/L_{\\rm Edd}$ is comparable to that inferred for the well studied Seyfert~2 NGC~1068, although since the absorber is completely opaque even to hard X--rays, the central luminosity can be estimated only indirectly.  These two AGN are at the opposite end of the scale to NGC~4258, which radiates at $\\sim 10^{-4}$ $L_{\\rm Edd}$ or less (Lasota et al.\\ 1996).  The mass accretion rates inferred from those values of $L/L_{\\rm Edd}$ put strong constraints on possible underlying radiation mechanisms.  While the recently popular advective disk models can produce X--ray hot flows up to about 10\\% of the Eddington limit, these collapse at higher $L/L_{\\rm Edd}$: our data show that NGC~4945 lies perilously close to this limits.  \\bigskip  {\\bf ACKNOWLEDGEMENTS:} We thank the RXTE satellite team for scheduling the observations allowing the daily sampling, Tess Jaffe for her help with the RXTE data reduction via her indispensable script {\\tt rex}, and Dr. Julian Krolik for his helpful comments on the manuscript.  This project was partially supported by ITP/NSF grant PHY94-07194, NASA grants and contracts to University of Maryland and USRA, and the Polish KBN grant 2P03D01816."
    },
    "0002/astro-ph0002255_arXiv.txt": {
        "abstract": "During the observation period 1997, the nearby Blazar Mkn~501 showed extremely strong emission and high variability. We examine multiwavelength aspects of this event using radio, optical, soft and hard X-ray and TeV data. We concentrate on the medium-timescale variability of the broadband spectra, averaged over weekly intervals.  We confirm the previously found correlation between soft and hard X-ray emission and the emission at TeV energies, while the source shows only minor variability at radio and optical wavelengths. The non-linear correlation between hard X-ray and TeV fluxes is consistent with a simple analytic estimate based on an SSC model in which Klein-Nishina effects are important for the highest-energy electrons in the jet, and flux variations are caused by variations of the electron density and/or the spectral index of the electron injection spectrum.  The time-averaged spectra are fitted with a Synchrotron Self-Compton (SSC) dominated leptonic jet model, using the full Klein-Nishina cross section and following the self-consistent evolution of relativistic particles along the jet, accounting for $\\gamma\\gamma$ absorption and pair production within the source as well as due to the intergalactic infrared background radiation. The contribution from external inverse-Compton scattering is tightly constrained by the low maximum EGRET flux and found to be negligible at TeV energies. We find that high levels of the X-ray and TeV fluxes can be explained by a hardening of the energy spectra of electrons injected at the base of the jet, in remarkable contrast to the trend found for $\\gamma$-ray flares of the flat-spectrum radio quasar PKS~0528+134. ",
        "introduction": "The BL Lac object Mkn 501 is very close ($z = 0.0337$, Ulrich et al. \\cite{ulrich}) and has  been studied extensively at all wavelengths. Together with its sister object Mkn 421, it was among the two first BL~Lac objects with known radio (Colla et al. \\cite{colla}), X-ray (Schwartz et al. \\cite{schwartz}) and TeV gamma-ray (Quinn et al. \\cite{quinn}, Bradbury et al. \\cite{bradbury}) counterparts. Recently, it was also marginally detected at photon energies $> 100$~MeV by the EGRET instrument on board CGRO (Kataoka et al. \\cite{kataoka}).  During 1997,  the object was found to be in an extreme high state with a TeV flux on average 20 times higher than in 1996 (Breslin et al. \\cite{breslin}). The source  exhibited strong variability on timescales of days with a possible quasi-periodically varying component with a timescale of about 25 days (Kranich et al. \\cite{kranich}). To complete the list of reasons for excitement for the observers, Mkn~501 reached in some of its flares fluxes of more than 10$^{-10}$~cm$^{-2}$s$^{-1}$ (above 1.5~TeV) - the most intense TeV emission ever measured so far from any astronomical object. However, the shortest observed variability timescale (5 hours, Aharonian et al. \\cite{hegraparti}) was significantly longer than that observed for Mkn~421 (Gaidos et al. \\cite{gaidos}).  All available TeV observatories monitored the event for several months (Samuelson et al. \\cite{whipple}, Aharonian et al. \\cite{hegraparti},\\cite{hegrapartii}, Hayashida et al. \\cite{telarray}, Djannati-Ata\\\"{\\i} et al. \\cite{cat}). The most complete dataset was produced by the HEGRA Cherenkov Telescope ``1'' (CT1) (Aharonian et al., \\cite{hegrapartii}) which was even able to observe Mkn~501 under the presence of moonlight, however with reduced sensitivity, thereby filling many gaps in the lightcurve. This telescope also obtained the confirmatory observations in 1996.  Also from HEGRA come probably the most accurate spectral measurements in the TeV regime. They were carried out by the HEGRA system of (at the time 4) Cherenkov telescopes (CTS) and are largely concurrent with the CT1 measurements, however with less time coverage (Aharonian et al. \\cite{hegraparti}).  The origin of the TeV $\\gamma$-ray emission and the reasons for its variability are still essentially unknown. The most popular models explain the TeV emission as near-infrared to UV photons which have been upscattered via the inverse Compton effect by very high energy electrons which are known from radio, optical and X-ray observations to be present in the jets of BL Lac objects. Possible sources of the seed photons for Compton scattering are the synchrotron radiation produced within the jet by the same population of relativistic electrons (synchrotron self-Compton model, SSC; Marscher \\& Gear \\cite{marscher}; Maraschi, Celotti \\& Ghisellini \\cite{maraschi}; Bloom \\& Marscher \\cite{bloom}), or radiation from outside the jet (external inverse Compton model, EIC). This external radiation could be the quasi-thermal radiation field of an accretion disk surrounding a supermassive black hole which is generally believed to power the relativistic jets. The accretion disk radiation can enter the jet either directly (Dermer, Schlickeiser \\& Mastichiadis \\cite{dermer92}, Dermer \\& Schlickeiser \\cite{dermer93}) or after being rescattered by circumnuclear material (Sikora, Begelman \\& Rees \\cite{sikora}, Blandford \\& Levinson \\cite{blandford}, Dermer, Sturner \\& Schlickeiser \\cite{dermer97}). It is also possible that synchrotron radiation produced within the jet and reflected by circumnuclear debris is the dominant source of soft photons during flares (Ghisellini \\& Madau \\cite{ghisellini96}, Bednarek \\cite{bednarek}), although it has been shown that this process is unlikely to be efficient in the case of BL~Lac objects (B\\\"ottcher \\& Dermer \\cite{boettcher98}).  As an alternative, Mannheim (\\cite{mannheim93}) has suggested a hadronic model in which protons are the primarily accelerated particles in the jet and the $\\gamma$-ray emission is produced by secondary pions and electron-positron pairs produced in photopion and photopair production interactions of the ultrarelativistic protons in the jet with external radiation. The time-averaged broadband spectrum of Mkn~501 has been fitted using this model by Mannheim et al. (\\cite{mannheim96}, \\cite{mannheim98}). However, the attempts to explain the short and intermediate-term variability of blazars with hadronic jet models have only just started (Rachen \\& Mannheim \\cite{rachen99}). For this reason, we concentrate on leptonic jet models in this paper since the variability time scales predicted by these models are in good agreement with the observed intraday variability of blazars.  Recently, Fossati et al. (\\cite{fossati}) have compared the broadband spectra of different types of $\\gamma$-ray emitting AGN and suggested a continuous sequence FSRQs (flat spectrum radio quasars) $\\to$ LBLs (low-frequency peaked BL~Lac objects) $\\to$ HBLs (high-frequency peaked BL~Lac objects), characterized by decreasing bolometric luminosity, decreasing dominance of the total energy output in $\\gamma$-rays compared to the emission at lower frequencies and a shift of the peak frequencies of the synchrotron and the $\\gamma$-ray component towards higher frequencies. Recent modelling efforts of various blazar-type AGN have revealed that this sequence is consistent with a decreasing importance of external radiation as a source of soft photons for Compton scattering in the jet (Ghisellini et al. \\cite{ghisellini98}). This suggests that the extreme HBLs like Mkn~501 or Mkn~421 can be well fitted with strongly SSC-dominated jet models, as was shown, e. g., by Mastichiadis \\& Kirk (\\cite{mastichiadis97}) and Pian et al. (\\cite{pian}).  From the SSC model one expects a strong correlation between the X-ray and the TeV emission. And indeed, by comparing the daily TeV measurements with daily averages from the All Sky Monitor (ASM) on board the Rossi X-Ray Timing Explorer (RXTE), a significant correlation with a most probable time-lag of 0 days (no time-lag) was found (Aharonian et al. \\cite{hegraparti,hegrapartii}). Furthermore, a clear overall X-ray high state was visible in the 1997 ASM measurements which coincided with the TeV high state (see also figure \\ref{fig-lightcurve}). And, as discovered by BeppoSAX, the X-ray peak of the spectral energy distibution (SED) had shifted from 10-20 keV in 1996 to 100-200 keV in 1997 (Pian et al. \\cite{pian}), consistent with the assumption that the X-ray and TeV-$\\gamma$-ray flares are produced by a more powerful electron acceleration, shifting the high-energy cutoff of the electron distribution to higher energies and hardening the X-ray spectrum. This  observation was, however, only made during one flare in 1997. In this paper we give evidence that the hardening took place during all flares in 1997. An alternative explanation for the synchrotron spectral changes, in terms of a steadily-emitting helical-jet model, has been presented by Villata \\& Raiteri (\\cite{villata99}).  Apart from the short-term variability on timescales of hours to days, there is also a longer-term variability in Blazars which has so far been investigated mainly in the optical regime (see e.g. Katajainen et al. (\\cite{katajainen}) and references therein) and recently also in X-rays (e.g. Mc Hardy \\cite{mchardy}). This variability shows remarkable amplitudes (e.g. 4.7 mag for Mkn 421 in the optical) and in the case of OJ 287 there is even evidence for periodicity.  In this respect Mkn 501 is not yet very well explored. Mkn 501 has been observed in the TeV regime since its discovery in 1995. The source showed low emission close to the sensitivity limits until the onset of the 1997 high state. Since the duty cycle of the TeV observatories is only about 10 \\%, this does not prove the absence of strong short flares prior to this high state, but the increased average intensity of the source can be described as an increase in flaring probability from 1996 to 1997 by at least an order of magnitude. This description is especially appropriate since even during the high state, the source returns to quiescent (comparable with 1996) levels of emission for periods of up to a few days.  These transitions are seen irrespective of the presence of short, strong flares on time-scales of a day.  The last observations in 1996 were made in August and found Mkn 501 still quiescent while the first observations in 1997 were made in February and found the source already flaring. The transition from low to high flaring probability obviously took place on timescales shorter than half a year. The correlated X-ray data from the RXTE ASM confirm this and show in addition that the change has been a smooth process over several months (see e.g. figure 8 in Aharonian et al. \\cite{hegrapartii}).  In this article we explore the multi-wavelength variability of Mkn 501 over medium timescales, and try to relate its behaviour to the physical parameters of a leptonic jet model. For this purpose we construct weekly SEDs using the HEGRA CT1 flux data and HEGRA CT System spectral data together with data from longer wavelengths, namely radio, optical, soft X-ray and hard X-rays. See table \\ref{tab-instruments} for an enumeration of the instruments and energy ranges.  This paper may represent an important step toward understanding the origin of strong changes in the flaring probability of Mkn 501 and Blazars in general, but our data is clearly not sufficient to give a final answer to this question.  All data except that from  HEGRA and the RXTE ASM are published here for the first time. The BATSE data are especially valuable since they confine the intensity at the X-ray peak of the SED. ",
        "conclusions": "We have presented broadband spectra of the extreme HBL Mkn~501 during its high state in 1997, including radio, optical, soft and hard X-ray, and TeV $\\gamma$-ray observations. In this study we concentrated on the medium-timescale variability, using weekly averaged SEDs. We confirmed the strong correlation between the TeV $\\gamma$-ray flux and the hard X-ray flux. This correlation was found to be non-linear and could be fitted with a second-order polynomial, in agreement with the expectation of an SSC dominated leptonic jet model, if the flux variations are related to fluctuations of the electron density in the jet and/or the spectral index of the electron spectrum at the time of injection into the jet.  The weekly averaged SEDs were fitted with a leptonic jet model, strongly dominated by the SSC process. With a few exceptions, this model yielded acceptable fits to the observed broadband spectra. The observed spectral variability of Mkn~501 could be explained mainly by variations of the electron spectral index. No clear correlation between the maximum electron energy and the hard X-ray and TeV $\\gamma$-ray fluxes on the 1-week timescale was found, in contrast to the short-term variability of Mkn~501. Pian et al. (\\cite{pian}) have shown that the intraday variability of this object is most probably related to an increase of $\\gamma_2$, leading to pronounced flares in hard X-rays, most probably on the synchrotron cooling timescale which is most likely $\\sim$~a few hours and thus much shorter than the 1-week timescale considered in this paper. Our result indicates that such synchrotron flares are isolated events and are at most weakly correlated to the activity of the source on the 1-week timescale.  Our result that the flaring behaviour on intermediate timescales is consistent with a hardening of the electron spectrum is in contrast to the flaring characteristics observed in quasars. Recently, Mukherjee et al. (\\cite{mukherjee99}) have investigated all available broadband data on the very luminous flat-spectrum radio quasar (FSRQ) PKS~0528+134, and found that its flaring behaviour is consistent with an increasing contribution of the external inverse-Compton component during flares, possibly related to an increase in the bulk Lorentz factor. The fits to the SEDs of PKS~0528+134 required that the average energy of relativistic electrons in the jet shifts towards lower values during flares, in contrast to the results found for Mkn~501. As pointed out by B\\\"ottcher (\\cite{boettcher99}), this implies that the synchrotron peak is expected to shift towards lower frequencies during flares of FSRQs, while Mkn~501 and Mkn~421 show clear evidence for a shift of the synchrotron peak to higher frequencies.  We point out that in the present study the magnetic field along the jet and the bulk Lorentz factor of individual blobs were fixed, so that we cannot confidently rule out variations of the Doppler factor, accompanied by appropriate changes of the electron injection spectrum, as the flaring mechanism for Mkn 501. However, the very moderate variability at optical frequencies, as observed in Mkn 501, leads us to consider this flaring mechanism less likely in this object since it would require a peculiar conspiracy between the Doppler factor and the electron spectrum to keep the optical flux at an approximately constant level."
    },
    "0002/astro-ph0002288_arXiv.txt": {
        "abstract": "\\rightskip = 0pt We have mapped the dense gas distribution and dynamics in the NW region of the Serpens molecular cloud in the CS(2--1) and \\n2hp(1--0) lines and 3~mm continuum using the FCRAO telescope and BIMA interferometer. 7 continuum sources are found. The \\n2hp\\ spectra are optically thin and fits to the 7 hyperfine components are used to determine the distribution of velocity dispersion. 8 cores, 2 with continuum sources, 6 without, lie at a local linewidth minimum and optical depth maximum. The CS spectra are optically thick and generally self-absorbed over the full 0.2~pc extent of the map. We use the line wings to trace outflows around at least 3, and possibly 4, of the continuum sources, and the asymmetry in the self-absorption as a diagnostic of relative motions between core centers and envelopes. The quiescent regions with low \\n2hp\\ linewidth tend to have more asymmetric CS spectra than the spectra around the continuum sources indicating higher infall speeds. These regions have typical sizes $\\sim 5000$~AU, linewidths $\\sim 0.5$~\\kms, and infall speeds $\\sim 0.05$~\\kms. The correlation of CS asymmetry with \\n2hp\\ velocity dispersion suggests that the inward flows of material that build up pre-protostellar cores are driven at least partly by a pressure gradient rather than by gravity alone. We discuss a scenario for core formation and eventual star forming collapse through the dissipation of turbulence. ",
        "introduction": "Stars form through the collapse of dense cores within molecular clouds. The detection and measurement of the motions associated with such star forming collapse appears to be secure (see reviews by Evans 2000; Myers, Evans, \\& Ohashi 2000). The focus has primarily been on isolated, individual star forming regions since these are the least complex cases to understand both observationally and theoretically. However, the majority of stars form in clusters (Zinnecker, McCaughrean, \\& Wilking 1993) so a broader understanding of star formation, including such fundamental issues as the origin of the IMF and the formation of massive stars, requires the study of how stars form in groups.  Relative to an individual star, the deeper potential well of a stellar cluster should imply faster, and possibly easier to measure, inward motions, but observations are complicated by the generally greater distance to clusters than to individual star forming regions studied so far, and by the overpowering luminosity of massive stars that limit the ability to image lower mass neighbors and thus to obtain a complete view of cluster birth. For example, observations of the massive cluster forming region, W49N, have produced evidence for a global collapse of material onto the cluster as a whole (Welch et al. 1987) and, at higher resolution, for infall onto bright individual protostars (Zhang \\& Ho 1997). However, it has not been possible to explore the formation of more moderate mass stars in these objects because of dynamic range limitations.  In this paper, we present millimeter line and continuum observations of a cluster forming region in the Serpens molecular cloud. This region is well suited for an exploration into the formation processes of stellar clusters because it is nearby (d=310~pc; de Lara, Chavarria-K, \\& Lopez-Molina 1991) and contains a moderately dense embedded cluster (stellar density $\\simgt 450$~pc\\eee; Eiroa \\& Casali 1992) with many highly embedded millimeter wavelength continuum sources (Testi \\& Sargent 1998) but no O or B stars. In addition, Williams \\& Myers (1997) and Mardones (1998) have found signatures of widespread infall motions in this region.  The combination of proximity and low mass makes it possible to identify individual star forming condensations and examine the structure and dynamics of the cluster on a core by core basis. With this goal, we mapped the cluster in the optically thick CS(2--1) and thin \\n2hp(1--0) lines with the FCRAO and BIMA telescopes. Their different optical depths allow us to probe cloud velocities from an outer envelope to the center (Leung \\& Brown 1977). The two species are also formed by different chemical pathways with abundances that depend on environment and time (Bergin et al. 1997) and thus their relative intensities offer information on the physical conditions and chemical age of the cores. In an earlier paper (Williams \\& Myers 1999a) we reported the discovery of a starless core that appears to be collapsing based on an analysis of a small part of the maps. Here, we report on the full dataset and investigate the dynamics of the dense gas across the whole cluster. We find a number of new cores, some starless, others with continuum sources, suggesting that new stars are continually being added to the group. There are spectroscopic signatures of outflow, infall, and localized dissipation of turbulence throughout the cloud which offer important clues about the physical processes involved in cluster formation. The observations are outlined in \\S2 and the data displayed in \\S3. An analysis of the data follows in \\S4, and we discuss our findings in \\S5, concluding in \\S6. ",
        "conclusions": "This paper presents millimeter wavelength continuum and spectral line observations of a young, embedded, low mass cluster forming region in the Serpens molecular cloud. 7 continuum sources are found at the resolution and sensitivity of these data. The distribution of these sources corresponds well with the \\n2hp\\ emission but the CS data presents a different appearance, with high velocity emission from outflows around 4 continuum sources, and central dips in spectral profiles from self-absorption (as shown on the large scale from singledish \\c34s\\ data and on the small scale from the \\n2hp). Away from the powerful outflow around S68N, the self-absorption is red-shifted which we interpret as indicating inward motions.  The \\n2hp\\ linewidth is dominated by non-thermal motions throughout the cluster. However, there are 8 regions where the non-thermal velocity dispersion reaches a local minimum. 6 are starless and 2 contain continuum sources.  They do not all coincide with peaks of the integrated intensity but they all stand out in maps of the peak temperature divided by the dispersion, a measure of the optical depth. The CS spectra toward these ``quiescent'' cores are particularly asymmetric, indicating relatively high infall speeds. Generally, the \\n2hp\\ dispersion increases and the CS blue-red temperature difference decreases with increasing distance from the core centers, and vice versa for the continuum sources. The correlation of CS asymmetry, related to infall speed, with \\n2hp\\ dispersion, related to the local turbulent pressure, suggests that the inward flows are at least partly pressure driven and that the cores formed through the localized dissipation of turbulence as envisioned by Nakano (1998) and Myers \\& Lazarian (1998). Such a scenario is consistent with the observed numbers of quiescent cores and Class 0 sources.  The singledish data alone shows a net inward motion onto the cluster. Although there is clearly considerable smearing of the detailed dynamics, this suggests that it may be fruitful to search for infall signatures in more distant clusters at $\\sim 0.1-0.2$~pc resolution. It will, of course, be important to observe other nearby cluster forming regions at higher resolution, $\\simlt 3000$~AU, to augment this study. Ultimately, the comparison of conditions in many clusters will give a clearer picture of their formation, show the effects of different environments, and, through an inventory of continuum sources and pre-protostellar cores, can be hoped to elucidate the origins of the stellar IMF (Motte et al. 1998; Testi \\& Sargent 1998)."
    },
    "0002/astro-ph0002241_arXiv.txt": {
        "abstract": "In the course of our ongoing search for serendipitous high-redshift \\lya\\ emission in deep archival Keck spectra, we discovered two very high equivalent width ($W_{\\lambda}^{\\rm obs} \\simgt 450$ \\AA, $2 \\sigma$) \\lya\\ emission line candidates at $z \\sim 3$ in a moderate dispersion ($\\lambda / \\Delta \\lambda \\simeq 1200$) spectrogram.  Both lines have low velocity dispersions ($\\sigma_v \\sim 60\\ \\kms$) and deconvolved radii $r \\approx 1\\ h_{50}^{-1}$ kpc.  We argue that the lines are \\lya, and are powered by stellar ionization.  The surface density of robust, high equivalent width \\lya\\ candidates is estimated to be $\\sim 3 \\pm 2\\ {\\rm arcmin}^{-2}$ per unit redshift at $z \\simeq 3$, consistent with the estimate of \\citet{Cowie:98}.  The \\lya\\ emission line source characteristics are consistent with the galaxies undergoing their first burst of star formation, \\ie\\ with being primeval.  Source sizes and velocity dispersions are comparable to the theoretical primeval galaxy model of \\citet{Lin:92} based on the inside-out, self-similar collapse of an isothermal sphere.  In this model, star formation among field galaxies is a protracted process.  Galaxies are thought to be able to display high equivalent widths for only the first $\\sim$ few $\\times 10^7$ yr.  This time is short in relation to the difference in look back times between $z=3$ and $z=4$, and implies that a substantial fraction of strong line-emitting galaxies at $z=3$ were formed at redshifts $z \\leq 4$.  We discuss the significance of high-equivalent width \\lya-emitting galaxies in terms of the emerging picture of the environment, and the specific characteristics of primeval galaxy formation at high redshift. ",
        "introduction": "Early theoretical models of primeval galaxies suggested that their \\lya\\ lines would be highly luminous and diffuse.  But searches on 4m-class telescopes confirmed no such emission lines \\citep{Pritchet:94, Thompson:95} and the search for galaxies at large look-back times turned to other strategies.  One highly successful technique is color-selection of the Lyman break in galaxies of redshift $z \\simeq 3$ \\citep{Steidel:92, Steidel:96a}.  Lyman-break galaxies (LBGs) are strongly star-forming, and compact relative to local ${\\cal L}^*$ galaxies \\citep{Giavalisco:96b}.  However, LBGs are not generally thought to be primeval, in part because they have lower \\lya\\ equivalent widths than expected of primeval galaxies, in part because interstellar metallic absorption lines (\\eg\\ \\siii, \\alii) are characteristic of most LBG spectra.  The theoretical upper limit equivalent width for \\lya\\ emission produced by stellar photoionization is thought to be $W_{\\lambda}^{\\rm rest} \\sim 100 - 200$ \\AA\\ \\citep{Charlot:93}, and can be expected to last only a few$\\times 10^7$ yr for a constant star formation rate (SFR).  Several examples of high-redshift, high equivalent width (non-AGN) \\lya\\ emission have been reported \\citep[\\eg][]{Dey:98, Weymann:98, Hu:99}.  Few secure examples of isolated, high equivalent widths that are not \\lya\\ exist in the literature \\citep[\\eg][]{Stockton:98, Stern:99b}.  About half of LBGs have \\lya\\ in emission \\citep{Steidel:96a, Steidel:98, Lowenthal:97}, with rest equivalent widths $\\sim 5 - 35$ \\AA.  Likewise, about half of local actively star-forming galaxies have \\lya\\ in emission, with rest equivalent widths in the range $W_{\\lambda} \\simeq 10 - 35$ \\AA\\ \\citep{Giavalisco:96a}, though a few have $W_{\\lambda} \\simgt 50 \\mbox{ \\AA}$.  Local studies show that whether \\lya\\ is seen in emission or absorption may depend heavily on the kinematics of the surrounding neutral \\hone\\ halo \\citep{Legrand:97, Kunth:98a} and the chance geometry of neutral gas and dust \\citep{Giavalisco:96a}.  However, $W_{\\lambda}^{\\rm rest}$ is shown to be weakly-correlated with metallicity \\citep{Terlevich:93, Giavalisco:96a}.  The selective absorption of resonantly scattered \\lya\\ photons, whose path length in the cloud is much longer than that of photons of modestly different wavelength, is thought to be the major factor in quenching the \\lya\\ line.  The partially extincted \\lya\\ line is characteristically asymmetric \\citep[\\eg][]{Dey:98, Kunth:98a, Legrand:97}, presumably caused by the absorption of photons from the blue side of the line by dust in an expanding shell, and compounded by the addition of red-shifted backscattering off the neutral shell expanding away from the observer \\citep[\\eg][]{Legrand:97}.  Though dust-free galaxies are the most likely to have high equivalent width \\lya\\ emission, it is clear that a more evolved galaxy can also have high equivalent width \\lya\\ emission emerging within specific zones of active star formation due to an incomplete covering factor.  Such emission may appear spatially offset from the region of star formation \\citep[\\eg][]{Bunker:00}.  Averaged over the whole galaxy, the continuum supplied by recent star formation with locally quenched \\lya\\ emission will reduce the integrated line equivalent width.  Deep spectra taken at the Keck telescopes regularly reveal serendipitous high-equivalent width, isolated emission lines \\citep[\\eg][]{Stern:99b}.  Indeed, the first confirmed galaxy at $z > 5$ was discovered in this manner \\citep{Dey:98}.  \\citet{Steidel:98} obtained a serendipitous spectrum of a strong \\lya\\ line with extremely faint continuum near the $\\langle z \\rangle \\simeq 3.09$ structure in the SSA22 field.  Strong emission lines can also be effectively targeted by using narrow band imaging, in conjunction with broad band imaging.  Recent observations, using narrow band imaging with a ``strong equivalent width'' ($W_\\lambda^{\\rm obs} \\geq 77$ \\AA) criterion \\citep{Cowie:98} and follow-up spectroscopy \\citep{Hu:98} have disclosed a population of galaxies with what is thought to be \\lya\\ in emission.  The galaxies tend to have very weak, sometimes undetectable, continua.  They are also are quite compact, as noted of the $z\\sim 2.4$ \\lya\\ emitters found in HST searches \\citep{Pascarelle:96, Pascarelle:98}.  \\citet{Cowie:98} estimate a surface density of 3.6 arcmin$^{-2}$ per unit redshift for flux densities $j > 2 \\times 10^{-17}~\\ergcm2s$, corresponding to line luminosities ${\\cal L}_{\\rm Ly\\alpha} >1.8 \\times 10^{42} \\ergs$ for an Einstein-de Sitter Universe with $H_0 = 50 \\kmsMpc$ at $z \\approx 3.4$.  We are conducting a search for serendipitous emission lines in deep archival Keck spectra.  Such a serendipitous search is an emission-line flux-limited survey, and may find emission lines over a wide range of redshifts.  Narrow-band surveys are sensitive to candidates over only small redshift ranges, and require follow-up spectroscopy to discriminate stellar \\lya\\ lines from metal and AGN emission lines. The identification of a population of primeval galaxies may provide important information about the epoch(s) of galaxy formation, including data relevant to the integrated global star formation rate, and the luminsity function of galaxies to photometric limits fainter than that accessible to color-selection surveys.  One drawback faced by all \\lya\\ emission line searches is possible confusion with other isolated, high equivalent width lines such as \\oii\\ and H${\\alpha}$.  Experience has shown that \\oii\\ is often mis-identified as \\lya.  When continuum is detected, we may use the Balmer series or the continuum depression to discriminate, and when absent we may search for another line, such as \\nv\\ or \\oiii. However, \\nv\\ is rarely detected in non-AGN, and the search for \\oiii\\ often fails because these lines would either fall out of the range of the spectrograph, or in the high-noise infrared part of the spectrum which, combined with the large intrinsic dispersion in \\oii/\\oiii\\ \\citep{Kennicutt:92}, can make an unambiguous verification untenable. Alternatively, one may use the known tendency of \\lya\\ lines to display a P-Cygni profile to identify this resonance line.  However, asymmetry cannot be a completely necessary criterion for \\lya\\ emission lines since some local sources apparently have symmetric, or ``pure'' profiles \\citep{Kunth:98b}.  When emission lines are unresolved, and have undetected continua, the last recourse is the knowledge that \\oii\\ rarely has an equivalent width greater than 100 \\AA\\ \\citep[but see][]{Stern:99b}.  For these reasons, unambiguous \\lya\\ lines are rare, and worthy of extra scrutiny.  We serendipitously found two isolated emission lines on a single slitlet from a mask centered on the SSA22 field, andq determined that they are in fact \\lya\\ lines --- examples of galaxies which are more compact and have higher \\lya\\ equivalent widths than known LBGs.  We report on their equivalent widths, intrinsic radii, velocity dispersions, and surface/volume density in order to place them in the picture of the evolving Universe emerging from studies of high-redshift galaxies.  The source targeted by the slitlet in which we found our emission line galaxies is the color-selected Lyman-break galaxy SSA22~C17; we analyze it in parallel for comparison. We adopt $H_0 = 50\\, h_{50}\\, \\kmsMpc, \\Omega_0 = 1,$ and $\\Lambda = 0$ unless otherwise stated.  At $z = 3$, for $\\Omega = 1 (0.1)$, 1\\farcs0 corresponds to $7.3\\, (12.3)\\, h_{50}^{-1}$ kpc. ",
        "conclusions": "We draw the following conclusions from this study:  \\begin{itemize}  \\item We suggest that the combination of high equivalent widths, low velocity dispersions, and small intrinsic sizes seen in the \\lya\\ lines of these two galaxies can be well-explained as their being extremely low-metallicity galaxies undergoing an initial burst of star formation--- hence, primeval galaxies.  These characteristics, markedly different from earlier expectations, suggests an inside-out, rather than a monolithic collapse formation mechanism.  \\item Our data, which imply that most galaxies at $z=3$ may have been formed at $z \\leq 4$, stand in apparent conflict with the $z=3$ and 4 UVLFs \\citep{Steidel:99}, which suggest that the comoving galactic density remained relatively constant during this interval. These may be reconciled, however, with a large merger rate that preserves the faint-end slope of the luminosity function.   \\item The agreement of the angular size, and the projected total luminosity of these emission line galaxies, with the predictions of the inside-out collapse of a nearly-isothermal \\hone\\ halo \\citep{Lin:92} suggests that the latter is a promising basis for modeling galaxy formation.  Our discussion has suggested that a rising star formation rate with a time-scale of $\\sim 5 \\times 10^7$ yr is plausible for many post-primeval galaxies.  \\item The isothermal collapse model requires that protogalactic clouds form in relatively isolated regions.  However, their collapse may be stimulated by the gradually increasing pressures of a protocluster environment, thus accounting for the apparent weak clustering on large scales noted of high-, as well as low-redshift \\lya\\ emission line galaxies.   \\end{itemize}"
    },
    "0002/astro-ph0002077_arXiv.txt": {
        "abstract": " ",
        "introduction": "The lateral distribution of the charged particle component of extensive air showers (EAS) carries information about the height of maximum of the EAS development. In NKG type parameterizations of the lateral distribution this information is associated with the so-called age parameter s, originally introduced by the analytic description of purely electromagnetic cascades for characterizing the actual stage of the EAS development. \\\\ In EAS experiments this parameter is usually extracted from fitting the distribution measured on observed level, assuming that this lateral parameter reflects the actual longitudinal EAS stage. Investigations of the parameter s have been performed on various altitudes, with the aim to gain information on the longitudinal EAS development and on the composition of primary cosmic rays [1-7]. For example, from the analysis of the zenith angle dependence of the average value of s it has been concluded that the mass composition gets either heavier primary energies larger than $10^{15}\\,$eV or the multiplicity of secondary particle production in hadronic interactions is unexpectedly increasing. \\\\ In the present contribution experimental age distributions, dependent on the zenith angle $\\theta$ of EAS incidence and of the shower size $N_e$ as extracted from an actual data set of the MAKET-ANI array, are communicated. As compared to earlier results \\cite{AH-avak} the statistical accuracy of the data is considerably improved thanks to various modernizations of the installation \\cite{AH-gagik}. The variation of the age parameter with the observation depth X is considered by a simplified approach. ",
        "conclusions": "The present results deduced from the data of the MAKET-ANI array are in good agreement with theoretical expectations. The analyses reveal that: \\\\ $\\bullet$ Average age parameter gradually decreases with increasing shower size from $10^5$ to $10^6$, and for $N_e>10^6$ it becomes almost constant.\\\\ $\\bullet$ The knee of \"young\" showers is sharper than knee of the all particle spectra.\\\\ $\\bullet$  The size spectra classified by different ages show different attenuation. \\\\ $\\bullet$  The change of age parameter with the zenith angle of EAS incidence can be related to the change of the EAS maximum with $N_e$.   \\vspace*{3.5ex} {\\noindent \\bf \\Large Acknowledgment \\\\} \\vspace*{-.5ex}  {\\it The report is based on scientific results of the ANI collaboration. The MAKET-ANI installation on Mt.Aragats has been setup as a collaboration project of the Yerevan Physics Institute (Armenia)  and the Lebedev Physics Institute (Moscow). The continuous contributions and assistance of the Russian colleagues in operating of the  installation and in the data analyses are gratefully acknowledged. In particular, we thank Prof. S. Nikolski and Dr. V. Romakhin for their encouraging interest to this work and useful discussions. \\\\ We would like to thank Prof. Dr. H. Rebel pointing on the importance of the elongation rate estimation, Dr. A. Haungs and Dr. Kh. Sanosyan for useful remarks. Corrections and suggestions made by Dr. S. Ostapchenko are highly appreciated. The assistance of the Maintenance Staff of the Aragats Cosmic Ray Observatory in operating the MAKET-ANI installation is highly appreciated. The work has been partly supported by the ISTC project A116.}"
    },
    "0002/astro-ph0002307_arXiv.txt": {
        "abstract": " ",
        "introduction": "A large number of spiral galaxies, including the Milky Way, carry magnetic fields \\cite{Kronberg}. With few exceptions, the galactic field strengths are measured to be a few times $10^{-6}$ G. This particular value has also been found at a redshift of $z=0.395$ \\cite{redshift} and between the galaxies in clusters.   Studies of the polarisation of synchrotron radiation emitted by galaxies with a face-on view, such as M51, reveal that their magnetic fields are aligned with spiral arms and density waves in the disk. A plausible explanation is that galactic magnetic fields were created by a mean-field dynamo mechanism \\cite{dynamo}, in which a much smaller seed field was exponentially amplified by the turbulent motion of ionised gas in conjunction with the differential rotation of the galaxy.  For the dynamo to work, the initial seed field must be correlated on a scale of $100$ pc, corresponding to the largest turbulent eddy \\cite{dynamo}. The required strength of the seed field is subject to large uncertainties; past authors have quoted $10^{-21\\pm 2}$ G as the lower bound at the time of completed galaxy formation. This would present a problem for most particle-physics and field-theory inspired mechanisms of magnetic field generation. However, in recent work with A.-C.\\ Davis and M.J.\\ Lilley \\cite{bbounds}, I have shown that the lower bound on the dynamo seed field can be significantly relaxed if the universe is flat with a cosmological constant, as is suggested by recent supernovae observations \\cite{supernovae}. In particular, for the same dynamo parameters that give a lower bound of $10^{-20}$ G for $\\Omega_0=1$, $\\Omega_\\Lambda=0$, we obtain $10^{-30}$ G for $\\Omega_0=0.2=1-\\Omega_\\Lambda$, implying that particle-physics mechanisms could still be viable. The observation at redshift $z=0.395$~\\cite{redshift} can also be accounted for with these parameters, but requires a seed field of at least $10^{-23}$ G \\cite{bbounds}.  The dynamo amplifies the magnetic field until its energy reaches equipartition with the kinetic energy of the ionised gas, $\\langle B^2/2\\rangle=\\langle\\rho v^2/2\\rangle$, when further growth is suppressed by dynamical back reaction. Thus a final field of $B_0\\approx 10^{-6}$ G results for any seed field of sufficient strength.  In order to explain the galactic field strength without a dynamo mechanism, one would require a strong primordial field of $10^{-3} (\\Omega_0 h^2)^{1/3}$ G  at the epoch of radiation decoupling $t_{\\rm dec}$, corresponding to a field strength $10^{-9} (\\Omega_0 h^2)^{1/3}$ G on comoving scales of $1$ Mpc. Future precision measurements of the CMB will put severe constraints on such a primordial field \\cite{CMB}. Moreover, magnetic fields on Mpc scales have been probed by observations of the Faraday rotation of polarised light from distant luminous sources, which give an upper bound of about $10^{-9}$ G \\cite{Kronberg}. The observation of micro-Gauss fields between galaxies in clusters presents an interesting dilemma. Because such regions are considerably less dense than galaxies, it is doubtful whether a dynamo could have been operative. Thus the intra-cluster magnetic fields, unless somehow ejected from galaxies, have formed directly from a primordial field stronger than $10^{-3} (\\Omega_0 h^2)^{1/3}$ G at $t_{\\rm dec}$. Such a field would certainly leave a signature in future CMB data \\cite{CMB}.  Particle-physics inspired models, which typically produce weak seed fields, lead to precise predictions and there have an advantage over astrophysical mechanisms, where the magnetic field strength is determined by complicated nonlinear dynamics, or solutions of general relativity with a magnetic field \\cite{Thorne}, where the field strength must be fixed by observations. With the possible exception of the last-mentioned model, there is no compelling scenario that produces a primordial field strong enough to eliminate the need for a dynamo.  Seed fields for the dynamo can be astrophysical or primordial. In the former category there is the important possibility that a seed field may arise spontaneously due to non-parallel gradients of pressure and charge density during the collapse of a protogalaxy \\cite{Kulsrud}. For the rest of this talk, however, I shall assume that the seed field is primordial. ",
        "conclusions": ""
    },
    "0002/astro-ph0002131_arXiv.txt": {
        "abstract": "Here we present the study of the relative angle between the accretion disk (or radio jet) and the galaxy disk for a sample of Seyfert galaxies selected from a mostly isotropic property, the 60$\\mu$m flux, and warm infrared colors. We used VLA A-array 3.6cm continuum data and ground based optical imaging, homogeneously observed and reduced to minimize selection effects. For parts of the analysis we enlarged the sample by including galaxies serendipitously selected from the literature. For each galaxy we have a pair of points ($i$,$\\delta$), which are the inclination of the galaxy relative to the line of sight and the angle between the jet projected into the plane of the sky and the host galaxy major axis, respectively. For some galaxies we also had information about which side of the minor axis is closer to Earth. This data is combined with a statistical technique, developed by us, to determine the distribution of $\\beta$ angles {\\it in 3 dimensions}, the angle between the jet and the host galaxy plane axis.  We found from an initial analysis of the data of the 60$\\mu$m sample, where Seyfert 1's and 2's were not differentiated, that the observed distribution of $i$ and $\\delta$ values can be well represented either by a homogeneous $\\sin\\beta$ distribution in the range $0^{\\circ}\\leq\\beta\\leq90^{\\circ}$, or $0^{\\circ}\\leq\\beta\\leq65^{\\circ}$, but not by an equatorial ring.  A more general model, which tested $\\beta-$distributions in the range $\\beta_1\\leq\\beta\\leq\\beta_2$, for different ranges of $\\beta_1$ and $\\beta_2$ values, required $\\beta_2$ to be larger than 65$^{\\circ}$ and gave preference for $\\beta_1$ smaller than 40$^{\\circ}$ - 50$^{\\circ}$.  An important result from our analysis was obtained when we distinguished if the jet was projected against the near or the far side of the galaxy, and differentiated between Seyfert 1's and Seyfert 2's, which showed that the model could not represent Seyfert 1's adequately. We found that the inclusion of viewing angle restrictions for Seyfert 1's, namely, that a galaxy can only be recognized as a Seyfert 1 if the angle between the jet and the line of sight ($|\\phi|$) is smaller than a given angle $\\phi_c$ and that the galaxy inclination $i$ is smaller than an angle $i_c$, gave rise to statistically acceptable models. This indication that there is a difference in viewing angle to the central engine between Seyfert 1's and Seyfert 2's is a direct and independent confirmation of the underlying concepts of the Unified Model.  We discuss possible explanations for the misalignment between the accretion disk and the host galaxy disk, which are:  warping of the accretion disk by self-irradiation instability, by the Bardeen-Petterson effect, or by a misaligned gravitational potential of a nuclear star cluster surrounding the black hole, as well as feeding of the accretion disk by a misaligned inflow of gas from minor mergers, capture of individual stars or gas from the nuclear star cluster, or the capture of individual molecular clouds from the host galaxy. ",
        "introduction": "We would expect, based on grounds of symmetry and simplicity, that the jets emanating from a Seyfert nucleus would emerge at right angles to the disk of the host spiral galaxy. The processes for bringing material close to the core (within the innermost 10pc) of a galaxy, either for the initial formation of the black hole or to provide fuel to the black hole, are of two sorts -- those that feed the nucleus from the visible gas reservoir in the galaxy disk, and those that feed the nucleus by introducing material from outside the galaxy. In the simplest pictures, fueling by both internal gas and by external gas favors co-alignment between accretion disk and galaxy disk, since most of the gas is in the galaxy disk and any gas added to it may rapidly end up there, either by shocks or by settling in the galaxy potential.  Since the jets are launched perpendicular to the accretion disk, the simplest assumption would be to see them aligned to the host galaxy minor axis.  However, the above scenario is flatly contradicted by the observations.  Investigations of the observed distribution of the angle $\\delta$, the difference between the position angle of the major axis of the galaxy and the position angle of the radio jet projected on the plane of the sky, shows that Seyfert galaxies can have jets along any direction, from aligned along the minor axis to aligned along the major axis (Ulvestad \\& Wilson 1984; Brindle et al. 1990; Baum et al. 1993; Schmitt et al. 1997; Nagar \\& Wilson 1999). It was shown by Clarke, Kinney \\& Pringle (1998), using data for a sample of Seyfert galaxies selected from the literature, that it is possible to obtain a reliable estimate of the distribution of the angle $\\beta$ {\\it in 3-dimensions} between the jet axis and the normal to the galaxy plane, by considering for each galaxy in the sample, the pair of values of $i$ and $\\delta$ ($i$ is the inclination of the galaxy to the line of sight). They conclude that the directions of the radio jets are consistent with being completely uncorrelated with the planes of the host galaxies (see also Nagar \\& Wilson 1999).  The observed random alignment between accretion disks and galaxy disks is intriguing, and the study of this effect is important for the understanding of the inner workings of Seyfert galaxies. This result may imply, for example, that recently suggested ideas about radiation induced warping of accretion disks (Pringle 1996, 1997, and Maloney, Begelman \\& Pringle 1996) come into play both to determine the directionality of the accretion disk and to produce the ionization cones, or that the warping is caused by the rapid spinning of the central black hole with spin misaligned with the spin of the galaxy (Bardeen-Petterson effect). Alternatively, this misalignment may result by the feeding of the accretion disk from gas outside the galaxy, by mergers, for example.  Here we study the distribution of angles $\\beta$ in a well defined sample of Seyfert galaxies, selected from a mostly isotropic property, the flux at 60$\\mu$m, and warm infrared colors (de Grijp et al. 1987, 1992). Most of the previous works in this subject used samples selected from the literature and were likely to suffer from selection effects. The example of one problem possibly caused by selection effects was the ``zone of avoidance'' found by Schmitt et al. (1997), a region of 20$^{\\circ}$ around the host galaxy minor axis where no jets were detected, but which was shown later by Nagar \\& Wilson (1999) to be due to the sample. Another problem of previous papers was the use of data collected from the literature, measured by different authors, using different methods, which resulted in uncertainties that cannot be quantified. We have addressed this problem by using radio and optical data observed, reduced and measured in a homogeneous way.   This paper is organized in the following way. In Section 2 we discuss the details about the data and the samples. In Section 3 we present the geometry of the problem and the statistical technique used to determine the distribution of $\\beta$ angles. In Section 4 we compare the data with the models, determine which $\\beta$-distribution and what kind of restrictions are required to better represent the observed data. In Section 5 we discuss a series of possible explanations for the observed misalignment between the accretion disk and the host galaxy disk, and finally in Section 6 we give a summary of the work. ",
        "conclusions": "In this paper we presented the study of the relative angle between the accretion disk and the galaxy disk for active galaxies hosted by spirals (Seyferts), using a sample selected by a mostly isotropic property, the flux at 60$\\mu$m, and warm infrared colors. This sample consists of 88 galaxies (29 Seyfert 1's and 59 Seyfert 2's), 33 of which show extended radio emission and are not in interacting systems (8 Seyfert 1's and 26 Seyfert 2's). Our study used VLA 3.6cm data taken by us, archival VLA data, ground based B and I images of the galaxy disks, as well as long slit spectroscopy. All the data were observed, reduced and analyzed in a similar way, to ensure a homogeneous dataset and minimize, as much as possible, selection effects. For parts of the analysis we also used an enlarged sample, which includes the 33 galaxies from the 60$\\mu$m sample, plus 36 serendipitous Seyfert galaxies selected from the literature, giving a total of 69 galaxies (20 Seyfert 1's and 50 Seyfert 2's). Most of the data for the serendipitous sample was obtained from the literature and not measured homogeneously as for the 60$\\mu$m sample.  For each galaxy we had a pair of measurements ($i$,$\\delta$), where $i$ is the galaxy inclination relative to the line of sight and $\\delta$ is the angle between the jet projected into the plane of the sky and the host galaxy major axis. For some of the objects, we also had the information about which side of the galaxy is closer to Earth, obtained through the inspection of dust lanes or from the rotation curve, assuming that the spiral arms are trailing. This information was used together with a statistical technique developed by us, to determine the distribution of angles $\\beta$, the angle between the jet and the host galaxy rotation axis. This technique tests different $\\beta-$distributions in the range $\\beta_1\\leq\\beta\\leq\\beta_2$, to determine which range of parameters $\\beta_1$ and $\\beta_2$ produces the most acceptable models.  From the assumption that the of a homogeneous $\\sin\\beta$ distribution in the range $0^{\\circ}\\leq\\beta\\leq90^{\\circ}$ and not differentiating between Seyfert 1's and 2's, we showed that the observed data and the models agree at the 67\\% level for the 60$\\mu$m sample and 81\\% level for the total sample (60$\\mu$m plus serendipitous sources). Using only a polecap ($0^{\\circ}\\leq\\beta\\leq65^{\\circ}$) we showed that the model and 60$\\mu$m data agree at the 69\\% level, while for an equatorial ring ($\\beta_1=\\beta_2=90^{\\circ}$) there is a bad agreement, only at the 0.8\\% level. Using a more general model, we tested which range of values $\\beta_1$ and $\\beta_2$ (where $\\beta_1\\leq\\beta_2$) are acceptable.  This showed that, independent of the sample (60$\\mu$m or total), $\\beta_2$ has to be larger than 65$^{\\circ}$-75$^{\\circ}$, and the acceptability of the models does not depend strongly on this value, but for $\\beta_1$ the agreement is better for values smaller than 40$^{\\circ}$-50$^{\\circ}$. The addition of the information about which side of the galaxy is closer to Earth, and if the jet is projected against the near or the far side of the galaxy, shows that the acceptable parameter range does not change considerably, but the maximum acceptability of the models was reduced. As discussed in Section 4.2, this result depends on the assumption that the dominant jet lies above the galaxy plane (as seen from Earth). We plan to test this hypothesis observing HI and free-free absorption against the weaker side of the radio jets.  An important result from our analysis appeared when we distinguished between Seyfert 1's and Seyfert 2's. The Seyfert 2's still have a good agreement with the models, and all the permitted $\\beta_1-\\beta_2$ parameter space is accepted at the 2$\\sigma$ level or higher. However, when we consider only the Seyfert 1's, the agreement is poor, with the maximum acceptability being only 7\\%. In order to solve this problem, we introduced a viewing angle restriction to the models, which is, a galaxy can only be recognized as a Seyfert 1 if the angle between the jet and the line of sight $|\\phi|$ is smaller than a given angle $\\phi_c$. We chose $\\phi_c=40^{\\circ}$ based on information from the literature, and show that this assumption increases the acceptability of the models.  This is, to our knowledge, the first time it is shown in a general way, using a statistically significantsample, that there is a difference in the viewing angle to the central engine of Seyfert 1's and Seyfert 2's, and is an independent confirmation of the Unified Model.   However, if we assume that all the Seyfert 2's have $|\\phi|>40^{\\circ}$ we find that NGC4388 would contradict this model, since the analysis of its data requires $|\\phi|\\leq21^{\\circ}$. We assumed, based on this and the analysis of the $i-$distribution of the Seyfert 1's in our sample, that a galaxy is only recognized as a Seyfert 1 if the jet is seen at an angle $|\\phi|\\leq\\phi_c$ and the host galaxy inclination is smaller than $i\\leq i_c$, otherwise it is a Seyfert 2. Comparing the models constructed assuming $\\phi_c=40^{\\circ}$ and $i_c=60^{\\circ}$ with the observed data for the 60$\\mu$m sample, we see that all permitted regions of the parameter space is acceptable, with a preference for small values of $\\beta_1$ and large values of $\\beta_2$.   As we discussed in the introduction, the simplest assumption suggests that the accretion disk is fed from gas in the galaxy disk, so we would expect that both disks have the same angular momentum vector. Since jets are supposed to be launched perpendicular to the accretion disk, the expectation would be to see all the jets aligned with the minor axis.  However, as shown above, this expectation is contradicted by our results, which clearly shows that the observed distribution of $\\delta$ and $i$ values can be represented by a homogeneous $\\beta-$distribution in the $0^{\\circ}\\leq\\beta\\leq90^{\\circ}$ range. We explored two main lines to explain the misalignment between the accretion disk axis and the host galaxy disk axis:  i) feeding of the accretion disk by aligned inflow from the galaxy disk, with the misalignment of the accretion disk; ii) feeding of the accretion disk by misaligned gas inflow.  In the case of aligned inflow, the randomness of the accretion disks could be due to warping of the accretion disk by self-irradiation instability, warping by the Bardeen-Petterson effect, or warping by a misaligned gravitational potential of a nuclear star cluster surrounding the black hole. In the case of misaligned inflow, the randomness of the jets could be due to misaligned minor mergers, capture of individual stars or gas from the nuclear star cluster, or the capture of individual molecular clouds from the host galaxy."
    },
    "0002/astro-ph0002182_arXiv.txt": {
        "abstract": "We discuss in  detail the information on large extra dimensions which can be derived in the framework of stellar evolution theory and observation. The main effect of large extra dimensions arises from the production  of the Kaluza-Klein (KK) excitations of the graviton. The KK-graviton and matter interactions are of gravitational strength, so the KK states never become thermalized and always freely escape. In this paper we first pay attention to the sun. Production of KK gravitons is incompatible with helioseismic constraints unless the $4+n$ dimensional Planck mass $M_s$ exceeds $300$ Gev/c$^2$. Next we show that stellar structures in their advanced phase of H burning evolution put much more severe constraints, $M_s > 3-4$  TeV/c$^2$, improving on current laboratory lower limits. ",
        "introduction": "\\label{first}  It is interesting to compare the energy loss due to KK-graviton production with the energy production from the pp chain at the center of the sun. The values of density, temperature and chemical composition derived from the SSM of \\cite{Bahcall98} are presented in Table \\ref{tabssm}. The results of other SSM calculations are similar, see e.g. \\cite{report}. Energy loss and production rates, computed according to the results of Appendix A, are compared in Table \\ref{tabepsi}. The most important contribution always comes from the photon-photon annihilation.  One expects that the solar solar structure would be drastically modified if the energy loss due to KK-gravitons becomes comparable with the nuclear energy production rate. In this way one can derive the following lower limits on $M_s$: \\begin{equation} \\label{ms2} n=2 \\,: \\quad M_s> 140 \\, GeV/c^2 \\end{equation} \\begin{equation} \\label{ms3} n=3 \\,: \\quad M_s>  3.5 \\, GeV/c^2 \\end{equation} This result which is essentially the same as that in ref. \\cite{Barger99} suggests that we concentrate on the $n=2$ case only. So far we assumed just a rough knowledge of the solar structure. One can expect that more detailed information, as that provided by helioseismology, provide more stringent constraints.  To understand in more detail the effect of KK-graviton production we have built solar models which include this additional energy loss. The energy generation subroutine was modified so as to include the KK-graviton loss and the stellar evolution code FRANEC \\cite{Ciacio} was run by varying the three free parameters of the model (initial helium abundance $Y_{in}$, initial metal abundance $Z_{in}$ and mixing length $\\alpha$) until it provides a solar structure (i.e. it reproduces the observed solar luminosity, radius and photospheric metal abundance  at the solar age).  As an example, we present here the case $M_s=0.2$ TeV/c$^2$. The main differences with respect to our SSM are depicted in Table \\ref{tab02} and Fig. \\ref{fig02}. Several features can be easily understood by observing that the solar model with KK-graviton production has to produce, now and in the past, a higher amount of nuclear energy, in order to compensate for the additional energy loss.  More hydrogen has been burnt into helium, and the initial helium abundance has to be reduced with respect to the SSM (otherwise one would get a stellar structure  which, being too much helium rich, would be presently overluminous). Consequently, the present photospheric helium abundance $Y_{ph}$ is decreased with respect to the SSM prediction.  Nevertheless, the central helium abundance is still somehow larger than in SSM and more energy is being produced in order to compensate for the KK-graviton losses. This is achieved with a somehow larger central temperature.  In the solar core, both temperature and ``mean molecular weight'' are thus higher than in the SSM, so that one cannot a priori decide for the behaviour of the sound speed. In fact Fig. \\ref{fig02} shows a decrease near the center and a significant increase near $R=0.2 R_\\odot$, i.e. in a region where helioseismic determinations are still very accurate.  These observations will be useful for determining the relevant observables which are sensitive to $M_s$ and which can be constrained by means of helioseismology.  In Fig. \\ref{figepsi} we also compare the nuclear energy production rate with the losses due to KK-gravitons.  The results are consistent with the qualitative energetics analysis discussed above. ",
        "conclusions": "We summarize here the main points of this paper:\\\\  \\noindent i)Helioseismic constraints on the sound speed in the energy production region and on the photospheric helium abundance rule out values of the $4+2$ dimensional Planck mass below $M_s= 0.3$ TeV/c$^2$.   \\noindent ii)The introduction of additional energy loss due to KK-graviton production cannot be a cure to the puzzles posed by SSM calculations. In particular, the predicted neutrino signals would be even larger than those of the SSM.  \\noindent iii)Observational constraints for Red Giant stars evolving in galactic globulars imply $M_s > 3-4$ TeV/c$^2$. This bound is stronger than that provided by accelerator, thus indicating how useful it is, and hopefully it will be, the synergetic use of terrestrial and stellar laboratories."
    },
    "0002/astro-ph0002461_arXiv.txt": {
        "abstract": "The TW Hydrae Association ($\\sim$55 pc from Earth) is the nearest known region of recent star formation.  Based primarily on the Hipparcos catalog, we have now identified a group of 9 or 10 co-moving star systems at a common distance ($\\sim$45 pc) from Earth that appear to comprise another, somewhat older, association (``the Tucanae Association'').  Together with ages and motions recently determined for some nearby field stars, the existence of the Tucanae and TW Hydrae Associations suggests that the Sun is now close to a region that was the site of substantial star formation only 10-40 million years ago.  The TW Hydrae Association represents a final chapter in the local star formation history. ",
        "introduction": "For many years, the only known star clusters within 70 pc of Earth were the rich Hyades (at $\\sim$ 45 pc) and the sparce U Ma nucleus (at $\\sim$25 pc).  Then a group of stars $\\sim$55 pc from Earth was established as a {\\it bona fide} association of T Tauri stars of age $\\sim$10 million years (Webb et al 1999; Sterzik et al 1999; Webb, Reid, \\& Zuckerman 2000).  This TW Hydrae Association (hereafter TWA) was unrecognized for many decades in spite of its being the nearest region of recent star formation (Kastner et al 1997). A ubiquitous signpost of newly formed stars has been a nearby molecular cloud.  But no interstellar cloud has been found near TW Hya despite multiple searches. Therefore, we ask: do additional unrealized young associations far from molecular clouds exist near Earth?   Different teams have used the Hipparcos catalog to search for previously unrecognized stellar associations. For example, Platais et al (1998) undertook ``A search for star clusters from the Hipparcos Data'' and list basic data for 5 ``very likely'' new clusters and associations as well as 15 ``possible'' ones. Yet, in their own words, ``At distances $<$100 pc the survey is incomplete as a result of the chosen search strategy.'' Of the 20 potential new groupings in their Table 1, the closest, which contains 11 members from Hipparcos, is 132 pc away.  By contrast, the ``Tucanae Association'' we propose in the present paper is only $\\sim$45 pc from Earth and much younger than the U Ma and Hyades clusters.  Just as the sparse U Ma cluster nucleus is accompanied by more numerous U Ma stream stars, we suggest that the Sun is embedded in a stream of stars with similar space motions (the ``Tucanae Stream''), with the Tucanae Association playing a role analogous to the U Ma nucleus.  Figures 1a and 1b depict the Tucanae Association and some stream stars, respectively.  These stars likely represent some of the younger, nearer, members of the more extensive Pleiades group or Local Association proposed by Eggen (e.g., see Jeffries 1995 and references therein).  \\begin{figure*} \\figurenum{1} \\epsscale{0.5} \\plotfiddle{tuc_f1a.ps}{1.5in}{90}{25}{25}{-70}{-28} \\plotfiddle{tuc_f1b.ps}{0in}{90}{25}{25}{95}{-7} \\plotfiddle{tuc_f1c.ps}{0in}{90}{25}{25}{260}{14}  \\caption{Potential moving group members in Table 1 (panels a. - The Tucanae Association or nucleus and b. - some of the Tucanae stream stars) and U Ma nucleus stars (panel c.) shown for comparison.  The R.A. and Dec axes are scaled to roughly show the true appearance on the sky though a distortion exists due to the square grid used.  The vectors represent the proper motion over 250,000 years. Probable members of the Tucanae Association and stream are indicated with large symbols and dark vectors; improbable members with small symbols and light vectors.  Two ``possible'' members are shown with large symbols and light vectors.  The bar represents 5 pc at the approximate distance to each group.  X-ray (RASS-BSC) sources are marked with squares, non-ray sources with triangles. Positions of four additional young stars with similar space motions which appear in the foreground of the Association are indicated by asterisks (see text).} \\end{figure*}   \\bigskip ",
        "conclusions": "We have found a previously unrecogized southern association - ``the Tucanae Association'' - which is only $\\sim$45 pc from Earth, but not quite as young as the recently established TW Hydrae Association.  Thus, in the past year, the number of known stellar associations within 60 pc of Earth has increased from two (the Hyades and U Ma) to four.  The existence of the Tucanae and TW Hydrae Associations resolves the mystery of how the Beta Pictoris moving group can be so young, 20 $\\pm$ 10 Myrs (Barrado y Navascues et al 1999) and yet so near to Earth (within 20 pc).  That is, 10-40 million years ago, the region through which the Sun is now passing experienced a significant era of star formation which produced the Beta Pictoris group, the two southern associations, and related stream stars.  We thank the crew at CTIO including N. Suntzeff, T. Ingerson, M. Fernandez, A. Gomez, E. Cosgrove \\& R. Venegas for critical help acquiring these data.  A special thanks to M. Smith for his flexibility with telescope maintenence schedules.  We are grateful to M.  Jura for calling the Mannings and Barlow paper to our attention and for pointing out the likelihood of proper motion companions to some of the stars in Table 1.  We thank Dr. Jura and E. E. Becklin and P. Lowrance for comments on drafts of the paper. This research was supported by NSF Grant 9417156 to UCLA and by the UCLA Center for Astrobiology."
    },
    "0002/astro-ph0002527_arXiv.txt": {
        "abstract": "We discuss observations  of the prompt X-- and $\\gamma$--ray emission and X--ray afterglow from GRB981226. This event has the weakest gamma-ray peak flux detected with the \\S\\ Gamma-Ray Burst Monitor. It shows an isolated X-ray precursor and the highest X-ray to gamma-ray fluence ratio measured thus far with the \\S\\ Wide Field Cameras. The event was followed up with the \\S\\ Narrow Field Instruments, and the X-ray afterglow was detected up to 10 keV. The afterglow flux is observed to rise from a level below the sensitivity of the MECS/LECS telescopes up to a peak flux of (5$\\pm 1) \\times 10^{-13}\\, \\ergcms$ in the 2-10~keV energy band. This rise is followed by a  decline according to a power law with index of 1.31$^{+0.44}_{-0.39}$. We discuss these results in the light  of the current GRB models. ",
        "introduction": "Follow-up observations of arcminute positions of Gamma-Ray Bursts (GRBs) provided by provided by the Wide Field Cameras on the \\S\\ satellite \\cite{Jager97,Boella97a} have shown that in most cases a fading X-ray counterpart identified as the GRB X-ray afterglow is detected \\cite{Frontera98}. The fading law  is generally  a smooth power law (e.g., Frontera et~al. 1999\\nocite{Frontera99}) except in two cases: GRB970508 \\cite{Piro98}, in which a late-time outburst of about 10$^5$~s duration started about 6$\\times 10^4$s after the main event, and GRB970828 \\cite{Yoshida99}, in which a peak structure of 4000~s duration appeared 1.25 $\\times 10^5$~s after the main event. Generally the afterglow light curves extrapolated back to the time of the bursts, are in agreement with the tail of the GRB time profiles \\cite{Costa97,Piro98,Frontera99}. This fact is considered as evidence that late afterglow emission and tail of the prompt GRB emission have the same origin \\cite{Frontera99}. In the fireball model scenario (see, e.g., the recent review by Piran 1999 \\nocite{Piran99}), this means that both the tail of the prompt emission and the late afterglow emission can be due to an external shock produced by the interaction of a relativistically expanding fireball with the Interstellar Medium (ISM). \\\\ Of the promptly localized GRBs  80\\% show X-ray afterglow, about 50\\% exhibit also optical emission, and 30\\% show radio emission. Radio emission is generally accompanied by optical emission, except in two cases: GRB990506 \\cite{Taylor99} from which also X-ray afterglow emission was observed with the RXTE/PCA experiment (Bacodine trigger 7549, Hurley 1999 \\nocite{Hurley99}), and GRB981226. For this event, in spite of several attempts to find the optical counterpart \\cite{Galama98,Rhoads98,Bloom98,Schaefer98,Castro-Tirado98,Wozniak98,Wozniak98,Lindgren99} none was identified. The best upper limit (R$\\sim$23~mag) to the optical flux of the GRB counterpart was reported by Lindgren \\etal (1999) \\nocite{Lindgren99}. Frail \\etal (1999a) \\nocite{Frail99a} reported the detection of a radio counterpart (VLA 232937.2--235553). The radio source peaked to 173$\\pm27$~$\\mu$Jy at 8.46~GHz at about 10 days after the burst. This time delay is tipical of all previously studied radio afterglows. However, the source declined relatively fast, following a power law decay ($\\propto t^{-\\delta_R}$) with $\\delta_R \\, = \\, 2.0 \\pm 0.4$ \\cite{Frail99b}. Two interpretations of this rapid decay have been proposed by Frail \\etal (1999b) \\nocite{Frail99b}: it is either the consequence of a jetted GRB or the result of a fireball shock in an ambient medium with variable density. An optical galaxy (R~=~24.85 mag), consistent with the radio transient position, has been proposed as the host galaxy of GRB981226 \\cite{Frail99b}.  GRB981226 was also followed up with the \\S\\ Narrow Field Instruments. A transient X-ray source was observed, which was proposed as the X-ray afterglow of GRB 981226 \\cite{Frontera98}. Here we report the properties of this afterglow emission  along with those of the prompt X-- and gamma--ray emission. We will discuss these properties in the light of the current models of GRBs. ",
        "conclusions": "GRB981226 shows the lowest gamma-ray peak flux among the bursts localized thus far with \\S\\ . In the log~N--log~P distribution of the GRBs observed with BATSE \\cite{Paciesas99}, it is located near the faint end of the distribution. The burst is marked by three peculiarities, two of which have been seen for the first time:  \\begin{enumerate}  \\item Of the bursts localized by BeppoSAX, this burst is the richest in the X-ray band:  the X-ray to gamma-ray fluence ratio of $1.4\\pm 0.4$ is the highest of all SAX bursts \\cite{Frontera99}. The peak energy $E_p\\sim 60$ keV is softer than that measured in other \\S\\  bursts. \\item An isolated X-ray precursor occurs about 180 s before the main event (Figures 1--2). Onset of X-ray emission before the gamma-rays has been observed in other GRBs \\cite{Laros84,Murakami91,Zand99,Feroci99}, but only another isolated X-ray precursor, started about 25~s before the start of $\\gamma$--rays, has been reported thus far \\cite{Laros84}. \\item The X-ray afterglow light curve is peculiar. As can be seen from Figure 4 (top), the afterglow emission is undetectable during first 2 hours of the start of the NFI observations (2$\\sigma$ upper limit of $1.5 \\times 10^{-13} \\ergcms$ in 2-10~keV), after which it rises rapidly to (5$\\pm 1) \\times 10^{-13} \\ergcms$, then undergoes a decline in the typical power law fashion (index $\\delta \\,= \\, 1.31^{+0.44}_{-0.39}$). Thus at least during epoch 11--13 hrs the afterglow is undetectable. \\end{enumerate}  Item (2) is an uncommon feature of GRBs. Assuming the synchrotron shock model (e.g., Piran 1999 \\nocite{Piran99}), it implies a an initial fireball expansion Lorentz factor smaller than that found in other GRBs (e.g., Frontera et~al. 1999\\nocite{Frontera99}). \\\\ Item (3) is the most interesting and mysterious aspect of this burst. In all SAX observed bursts to date, the afterglow emission at X-ray wavelengths begins more or less as soon as the main gamma-ray event ends. Also in the case of GRB981226 we have evidence that the afterglow starts during the main event. Indeed the peak flux of the X-ray prompt emission ($\\gtrsim 300$~$\\mu$Jy) has the same order of magnitude as that (173$\\pm27$~$\\mu$Jy at 8.46~GHz) observed in the radio band $\\sim$10 days after the burst \\cite{Frail99b}. On the contrary, the X-ray  flux measured by us about 13~hrs after the main event is about 4 orders of magnitude lower than the radio peak flux. Now similar values of peak fluxes in X-rays and in the radio band  are expected in  quasi-adiabatic cooling shocks of relativistically expanding fireballs \\cite{Sari98}. Thus a simplest conclusion is that the afterglow emission began immediately after the burst (phase E in fig.~4), rather than 13~hrs later, and then it was strongly reduced or completely ceased, and eventually restarted 13 hours after the burst. If we accept this explanation, then we must explain the physical cause for this rather extended gap. \\\\ The afterglow emission is attributed to the forward shock, i.e. shock of the ambient gas particles swept up by the advancing blast wave. A cessation of X-ray afterglow would require that there be no ambient gas (or very reduced density) and the resurgence of the afterglow at 13 hr would then require increased density -- in short, a cavity surrounding the explosion. The rapid decline of the X-ray and radio afterglow require that the ambient density not be a constant but decrease with the radius \\cite{Chevalier99}.  Indeed, one expects such a circumburst medium around massive stars to have a complicated geometry. For example, a  star which exploded as a blue supergiant would have suffered two episodes of mass loss. During the first phase, the red supergiant phase, the wind speed is low leading to a rich circumstellar medium. During the next phase, the blue supergiant phase, the fast blue supergiant wind sweeps up the circumburst medium shaped by the red supergiant wind with the net result of a low-density cavity surrounded by a dense shell at the outskirts.  If one accepts this explanation then we have found an additional evidence linking GRBs to massive stars. In this scenario, we find a satisfying explanation also for item (1) and (2). The X-ray richness is because the blast wave picks up matter (baryons) as it tunnels through the envelope of the massive star. The resulting large baryon content leads to lower $E_p$ and hence more X-ray emission.  We attribute the percursor emission to emission from shock breakout that is a natural consequence of models in which GRBs arise from the death of massive stars \\cite{MacFadyen99}.  The rapid decline in the radio and X-ray afterglow is because of the radial density gradient in the circumburst medium. Finally, the absence of an optical afterglow could well be due to extinction towards the GRB. If this GRB arises from death of a massive star then most likely the progenitor was in a dusty region and hence the extinction.  We tend to exclude that the X-ray increase of the afterglow  observed $\\sim$13~hrs after the burst is a signature of a supernova emission. Given that the X-ray flux is similar to that measured for the X-ray counterpart of SN1998bw \\cite{Pian99}, one would expect a moderate distance for GRB981226 (z$\\lesssim$~0.01) and therefore a significant detection in the optical band also in the case of a large extinction."
    },
    "0002/astro-ph0002056_arXiv.txt": {
        "abstract": "Recent detailed simulations of Galactic Chemical Evolution have shown that the heavy elements, in particular [Fe/H], are expected to exhibit a weak, or absent, correlation with stellar ages in the early Galaxy due to the lack of efficient mixing of interstellar material enriched by individual Type II supernovae.  A promising alternative ``chronometer'' of stellar ages is suggested, based on the expectation that the light elements Be and B are formed primarily as spallation products of Galactic Cosmic Rays. ",
        "introduction": "It has become clear, from a number of lines of recent evidence, that the early evolution of the Galaxy is best thought of as a stochastic process.   Within the first 0.5-1 Gyr following the start of the star formation process, chemical enrichment does not operate within a well-mixed uniform environment, as was assumed in the simple one-zone models that were commonly used in past treatments of this problem.  Rather, the very first generations of stars are expected to have their abundances of heavy elements set by local conditions, which are likely to have been dominated by the yields from individual SNeII.  The seeds of this paradigm shift can be found in the observations, interpretations, and speculations of McWilliam et al. (1995), Audouze \\& Silk (1995), and Ryan, Norris, \\& Beers (1996).  Models which attempt to incorporate these ideas into a predictive formalism have been put forward by Tsujimoto, Shigeyama, \\& Yoshii (1999; hereafter TSY), and Argast et al.  (2000). Although they differ in the details of their implementation, and in a number of their assumptions, both of these models rely on the idea of enhanced star formation in the high-density shells of SN remnants, and the interaction of these shells of enriched material with a local ISM.  The predictions which result are similar as well: (1) Both models are capable of reproducing the observed distributions of abundance (e.g., [Fe/H]) for stars in the tail of the halo metallicity distribution function (Laird et al. 1988; Ryan \\& Norris 1991; Beers 1999), and (2) Both models predict that the abundances of heavy elements, such as Fe, are not expected to show strong correlations with the ages of the first stars, at least up until an enrichment level on the order of [Fe/H]$ \\sim -2.0$ is reached, i.e., at the time when mixing on a Galactic scale is possible (roughly 1 Gyr following the initiation of star formation).  Suzuki, Yoshii, \\& Kajino (1999; hereafter SYK, see also Suzuki, Yoshii, \\& Kajino, this volume) have extended the SN-induced chemical evolution model of TSY to include predictions of the evolution of the light element species $^9$Be, $^{10}$B, and $^{11}$B, based on secondary processes involving spallative reactions with Galactic Cosmic Rays (hereafter GCRs).  Recently, Suzuki, Yoshii, \\& Beers (2000) have considered the extension of this model to the prediction of $^6$Li and $^7$ Li, and demonstrate that they naturally reproduce the recently detected slope in the abundance of Li in extremely metal-poor stars noted by Ryan, Norris, \\& Beers (1999; see also Ryan this volume).  It is particularly encouraging that the same stochastic star-formation models which reproduce the observed trends of some (but not all) heavy elements, such as Eu, Fe, etc., also obtain predictions of the light element abundance distributions that match the available observations quite well, with a minimum of parameter tweaking.  In this contribution we summarize one of the more interesting predictions of the TSY/SYK class of models, that the abundances of the light elements Be and B (hereafter, BeB) might be useful as stellar chronometers in the early Galaxy (a time when the heavy element ``age-metallicity'' relationships are not operating due to the lack of global mixing).  It appears possible that, with refinement of the modeling, and adequate testing, observations of BeB for metal-poor stars may provide a chronometer with ``time resolution'' on scales of tens of Myrs. ",
        "conclusions": ""
    },
    "0002/astro-ph0002260_arXiv.txt": {
        "abstract": "The development of sensitive large format imaging arrays for the infrared promises to provide revolutionary capabilities for space astronomy. For example, the Infrared Array Camera (IRAC) on SIRTF will use four $256\\times 256$ arrays to provide background limited high spatial resolution images of the sky in the 3 to 8 $\\mu$m spectral region. In order to reach the performance limits possible with this generation of sensitive detectors, calibration procedures must be developed so that uncertainties in detector calibration will always be dominated by photon statistics from the dark sky as a major system noise source. In the near infrared, where the faint extragalactic sky is observed through the scattered and reemitted zodiacal light from our solar system, calibration is particularly important. Faint sources must be detected on this brighter local foreground.  We present a procedure for calibrating imaging systems and analyzing such data. In our approach, by proper choice of observing strategy, information about detector parameters is encoded in the sky measurements. Proper analysis allows us to simultaneously solve for sky brightness and detector parameters, and provides accurate formal error estimates.  This approach allows us to extract the calibration from the observations themselves; little or no additional information is necessary to allow full interpretation of the data. Further, this approach allows refinement and verification of detector parameters during the mission, and thus does not depend on {\\it a priori} knowledge of the system or ground calibration for interpretation of images. ",
        "introduction": "The Infrared Array Camera (IRAC) (Fazio \\etal\\ 1998) will employ four $256\\times256$ imaging infrared arrays and the cooled telescope of the SIRTF to produce images of the sky which are limited by the photon statistics from the natural background, which, in this spectral region (8-25 $\\mu$m), is dominated by scattered and emitted light from the zodiacal dust particles. This will be typical of future applications of infrared detectors in space. In order to produce high quality images in the presence of this strong background, the relative response of the different pixels in the detector array must be known to high precision. A technique must be developed that allows the detector properties to be determined in operation, so that the requisite stability can be experimentally verified, and changes in response can be measured and included in the analysis of the data. We present a technique by which the detector properties are determined simultaneously with the estimates of sky brightness, and formal errors developed for both instrument and sky parameters.  We observe the same area of the sky with the detector array at a number of spatially offset positions. These observations are used to set up a system of linear equations involving both sky brightness and detector properties. In solving this system of equations, we can deduce the sky brightness and detector gain and offset parameters. By appropriate choices of offset spacings and sky brightness distributions, this technique allows us to continuously improve our knowledge of the detector properties or detect changes. This approach embeds the relative calibration of the detector array into the survey process; all information required to produce an internally consistent survey can deduced from the survey itself. Since the data on which the calibration is based is the survey itself, it is the way to calibrate the data which is, in some sense, least susceptible to systematic errors. In the case that an {\\it a priori} calibration is used, this technique offers a method to test internal consistency.  In this paper, we describe this least squares solution for sky and detector properties, and suggest implementations of the technique for the IRAC instrument. We present the analysis of synthetic Wide-Field Infrared Explorer (WIRE) data and real Hubble NICMOS data, in which we derive the sky brightness, detector gain and detector offset. (We had planned a demonstration of the technique on the Wide-field Infrared Explorer, but its unfortunate demise renders the point moot.) The results are encouraging, and form the basis of our plans for the analysis of the IRAC imaging data. Optimization of the observational strategy to produce the best encoding of the detector parameters in the survey observations is treated in a separate paper (Arendt, Fixsen \\& Moseley 2000). This approach can offer significant insurance to the observer, in that regardless of the availability or applicability of independent relative calibration data for the instrument, sufficient information is present in the observations themselves to allow the relative calibration of the data. This provides the capability for the observer to validate the statistical properties of the data or to calibrate it as required.  Least squares techniques are an important staple of model fitting. In this paper, we use a least squares technique, combined with sampling over a wide range of spatial scales, to produce an intensity calibration for the imaging system. Investigators have long used \"sky flats\" to produce estimates of system response (e.g. Joyce, 1992). In this process, images are taken at a variety of positions around the object of interest. These images are often processed using median filtering to produce estimates of detector response. In this paper , we derive the full algorithm for optimal use of the sampled data for intensity calibration of an imaging detector. This algorithm then allows us to ask more sophisticated questions important for planning observations, such as comparing the relative goodness of different sampling procedures (Arendt et al. 1999). This algorithm provides the optimal tool for calibrating imaging detectors; if the algorithm does not produce reliable results, it is indicative of incompleteness in the sampling of the sky. The algorithm provides an optimal detector calibration based on the data provided it. If a priori information about the detector is known, the algorithm can be adjusted to include it.  Other algorithms have been described in the literature for analyzing dithered image data. The drizzle method (Fruchter and Hook, 1998) is an approach for combining undersampled dithered images to produce a single combined image with improved resolution and signal-to-noise ratio. However, this technique is a means of a producing final image from calibrated data, and is not intended as a method of deriving the detector calibration.  Future observatories will generate survey data. The accuracy of analysis of these data will depend on a clear understanding of the statistical properties of the uncertainties in the data, their level, and spatial and temporal correlations. We present an approach for the analysis of such data, with specific application to the imaging data from the SIRTF IRAC instrument.  This comprehensive least squares approach has been successfully applied to the analysis of the data from the FIRAS instrument on COBE, in which a complex instrument model was required (Fixsen \\etal 1994). ",
        "conclusions": "As demonstrated with the simulated WIRE data the program can calculate the gains and offsets to the theoretical limit on the accuracy if it is given a good model of the data. As shown with the NICMOS data the program works reasonably well on real data as well even with he normal complexities of real errors and uncertainties. The uncertainties are calculated, and the correlations can be calculated with minor changes to the program. These allow the user to interpret the result without ad hoc assumptions or guesses about how the errors are related. The speed of the program allows modest data sets to be processed in a few minutes and with the availability of machines with large memories will allow the large data sets of the future to be processed in reasonable times."
    },
    "0002/astro-ph0002110_arXiv.txt": {
        "abstract": "{\\sc Simbad} is the reference database for identification and bibliography of astronomical objects. It contains identifications, `basic data', bibliography, and selected observational measurements for several million astronomical objects.  {\\sc Simbad} is developed and maintained by CDS, Strasbourg. Building the database contents is achieved with the help of several contributing institutes. Scanning the bibliography is the result of the collaboration of CDS with bibliographers in Observatoire de Paris (DASGAL), Institut d'Astrophysique de Paris, and Observatoire de Bordeaux.   When selecting catalogues and tables for inclusion, priority is given to optimal multi-wavelength coverage  of the database, and to support of research developments linked to large projects. In parallel, the systematic scanning of the bibliography reflects the diversity and general trends of astronomical research.  A  WWW interface to {\\sc Simbad} is available at:  {http://simbad.u-strasbg.fr/Simbad}. ",
        "introduction": "\\subsection{The CDS} The Centre de Donn\\'ees astronomiques de Strasbourg (CDS) defines, develops, and maintains services to help the astronomers find the information they need from the very rapidly increasing wealth of astronomical information, and particularly of on-line information.  CDS is operated at the Strasbourg astronomical Observatory, under an agreement between French Institut National des Sciences de l'Univers (INSU) and Universit\\'e Louis Pasteur, Strasbourg (ULP). CDS personnel created and implemented the {\\sc Simbad}  data bank and maintain its data and software system.  A detailed description of the CDS on-line services can be found, e.g., in Egret et al.\\ (\\cite{cds-amp2}) and in Genova et al.\\ (\\cite{cds-hub}, \\cite{cds}, \\cite{cds2000}), or at the CDS web site\\footnote{{\\em Internet address:} http://cdsweb.u-strasbg.fr/}. Questions or comments about the CDS services can be sent to the hot line {\\it question@simbad.u-strasbg.fr}.  \\subsection{SIMBAD}  The {\\sc Simbad} database contains identifications, `basic data', bibliographical references, and selected observational measurements for more than 2.7 million astronomical objects (November 1999). Data and information published in {\\sc Simbad} come from selected catalogues and tables and from the whole astronomical literature.  The specificity of the {\\sc Simbad} database is to organize the information per astronomical object, thus offering a unique perspective on astronomical data. This is done through a careful cross-identification of objects from catalogues, lists, and journal articles. The ability to gather together any sort of published observational data  related to stars or galaxies has made {\\sc Simbad} a key tool used worldwide for all kinds of astronomical studies.   {\\sc Simbad} is the acronym for  {\\sl S}et of {\\sl I}dentifications, {\\sl M}easurements and             {\\sl B}ibliography for             {\\sl A}stronomical {\\sl D}ata.  The main access point to {\\sc Simbad} is the WWW home page\\footnote{http://simbad.u-strasbg.fr/Simbad}; there is a mirror copy at SAO, Harvard\\footnote{http://simbad.harvard.edu/Simbad}.  \\subsection{Historical background}  Building a reference database for stars  -- and, later, for extragalactic objects and all astronomical objects outside the Solar System -- has been the first goal of the CDS: {\\sc Simbad} is the result of an on-going effort which started soon after the creation of CDS in 1972. {\\sc Simbad} was created by merging the Catalog of Stellar Identifications (CSI, Ochsenbein et al.\\ \\cite{csi}) and the Bibliographic Star Index (Ochsenbein \\cite{bsi}) as they existed until 1979.  The resulting data base (at that time, about 400,000 objects, mainly stars) was then expanded by the addition of source data from the catalogs and tables, and by  new literature references. The database was extended to galaxies and other non-stellar objects in 1983 (Dubois et al. \\cite{dubois83}). For details about the early developments of {\\sc Simbad} see Egret (\\cite{story}).  The first on-line interactive version of {\\sc Simbad} was released in 1981, and operated at  the Strasbourg Cronenbourg computer center until December 1984, when it was moved to Universit\\'e Paris-Sud at Orsay, and operated there until June 30, 1990. The database is now hosted on a Unix server, at the Strasbourg Observatory.  The original command line interface has been complemented by an interactive X-Window interface ({\\sc XSimbad}) in 1994, and by a World-Wide Web interface in 1996. There is also a client/server mode,  providing quick responses to simple queries, essentially for the name resolution in archives and information systems (see Section~\\ref{interface}).  For descriptions of earlier stages of the database, see Heck \\& Egret (\\cite{messenger}), and Egret et al.\\ (\\cite{ampersand}). ",
        "conclusions": ""
    },
    "0002/astro-ph0002326_arXiv.txt": {
        "abstract": "We present first preliminary data from the L3+Cosmics experiment and results from Monte Carlo simulations of multi-muon events as observed 30 m underground. ",
        "introduction": "The muon component of extensive air showers (EAS), due to the long muon range in the Earth's atmosphere, carries a wealth of information about the shower development. Study of multi-muon events gives an insight into the primary cosmic ray composition and the physics of high energy hadronic interactions. The L3 detector, situated 30 m underground, offers interesting possibilities to detect and study such events~\\cite{PL97}, which are complementary to the data collected in traditional cosmic ray experiments. The hadron component of EAS is absorbed, while the muon component is detected with low threshold (typically, if we exclude access shafts, 15 GeV) and high momentum and spatial resolution by the sophisticated tracking system of the L3 detector. The muon spectrum can be measured up to 2 TeV with high precision. The multi-muon event rate is high enough to make studies of the knee region possible with one year of data taking.  This year, 5 billion triggers were collected with the full L3+Cosmics setup. The independent readout and data acquisition system allows us to take data in parallel with L3. The acceptance of the setup is 200~$\\rm m^2 sr$. The angular resolution is better than 3.5 mrad for muons above 100 GeV and zenith angles from 0 to $\\rm 50^{\\circ}$. The momentum resolution is 5.0 \\% at 45 GeV. It is calibrated with $\\rm Z \\rightarrow \\mu^+\\mu^-$ events from the LEP calibration runs, where the muon momentum is known exactly. ",
        "conclusions": ""
    },
    "0002/astro-ph0002489_arXiv.txt": {
        "abstract": "Close-in extrasolar giant planets may be directly detectable by their reflected light, due to the proximity of the planet to the illuminating star. The spectrum of the system will contain a reflected light component that varies in amplitude and Doppler shift as the planet orbits the star.  Intensive searches for this effect have been carried out for only one extrasolar planet system, $\\tau$~Bo\\\"o. There exist several other attractive targets, including the transiting planet system HD~209458. ",
        "introduction": "Radial velocity surveys of nearby Sun-like stars have uncovered a population of close-in orbiting companions of roughly Jupiter mass.  The ten such objects with semi-major axes \\mbox{$a \\la 0.1$ AU} are listed in Table~1, along with the values for the semi-major axes, (minimum) masses, spectral types of the stars, and equilibrium temperatures (calculated from the estimated values for the semi-major axes, stellar radii and effective temperatures, and assuming a Bond albedo of $A$).  \\begin{table} \\caption{Close-In Planets Ordered by Increasing Semi-Major Axis} \\begin{tabular}{l c c c c} \\tableline Star & Spectral Type & $a$ (AU) & $M_p \\ (M_{\\rm Jup})$ & $T_{\\rm eq} (1-A)^{-\\frac{1}{4}}$ (K) \\\\ \\tableline HD 187123 & G3~V & 0.042 & 0.52/$\\sin i$ & 1400 \\\\ HD 75289 & G0~V & 0.046 & 0.42/$\\sin i$ & 1600 \\\\ $\\tau$ Bo\\\"o & F7~V & 0.046 & 3.87/$\\sin i$ & 1600 \\\\ HD 209458 & G0~V & 0.047 & 0.69 & 1500 \\\\ 51 Peg & G2~V & 0.051 & 0.45/$\\sin i$ & 1300 \\\\ $\\upsilon$ And & F8~V & 0.059 & 0.71/$\\sin i$ & 1500 \\\\ HD 217107 & G7~V & 0.072 & 1.28/$\\sin i$ & 1000 \\\\ HD 130322 & K0~V & 0.088 & 1.08/$\\sin i$ & 800 \\\\ 55 Cnc & G8~V & 0.11 & 0.84/$\\sin i$ & 700 \\\\ Gl 86 & K0~V & 0.11 & 3.6/$\\sin i$ & 700 \\\\ \\tableline \\tableline \\end{tabular} \\end{table}  A successful spectroscopic detection of an extrasolar planet in reflected light would yield the inclination, and hence the planetary mass, and would also measure a combination of the planetary radius and albedo. Furthermore, it would open the way to direct investigation of the spectrum of the planet itself.  Conversely, a low enough upper limit would provide useful constraints on the radius and albedo of the companion.  The predicted albedo of a close-in extrasolar giant planet has been the focus of recent theoretical work (Marley et al. 1999; Seager, Whitney, \\& Sasselov 2000), and the possible values range by several orders of magnitude based on the atmospheric temperature, chemical composition (see, e.g., Burrows \\& Sharp 1999), and the presence (or absence) of atmospheric condensates, and their respective size distributions. ",
        "conclusions": ""
    },
    "0002/astro-ph0002440_arXiv.txt": {
        "abstract": "Using both the Short- and Long-wavelength Spectrometers on board the Infrared Space Observatory (ISO), we have obtained infrared spectra of the Herbig Ae systems AB~Aur and HD~163296. In addition, we obtained ground-based $N$ band images of HD~163296. Our results can be summarized as follows: (1) The main dust components in AB~Aur are amorphous silicates, iron oxide and PAHs; (2) The circumstellar dust in HD~163296 consists of amorphous silicates, iron oxide, water ice and a small fraction of crystalline silicates; (3) The infrared fluxes of HD~163296 are dominated by solid state features; (4) The colour temperature of the underlying continuum is much cooler in HD~163296 than in AB~Aur, pointing to the existence of a population of very large (mm sized) dust grains in HD~163296; (5) The composition and degree of crystallization of circumstellar dust are poorly correlated with the age of the central star. The processes of crystallization and grain growth are also not necessarily coupled. This means that either the evolution of circumstellar dust in protoplanetary disks happens very rapidly (within a few Myr), or that this evolution is governed by factors other than stellar mass and age. ",
        "introduction": "Herbig Ae/Be stars are intermediate-mass pre-main sequence stars surrounded by disks of gas and dust which might be the site of on-going planet formation (see Waters \\& Waelkens 1998 for a recent review). AB~Aur (A0Ve) and HD~163296 (A3Ve) belong to the best studied Herbig stars and are sometimes considered prototypical for the entire class. As early as 1933, Merrill \\& Burwell remarked upon the similarity of both systems, which was confirmed by numerous subsequent authors. Apart from the fact that the stellar mass, effective temperature and age of AB~Aur and HD~163296 are nearly identical (van den Ancker et al. 1997, 1998), the similarity between the two systems also extends to their circumstellar environment: both AB~Aur and HD~163296 are examples of relatively isolated star formation, and are not hindered by confusion with other sources (Henning et al. 1998; Di Francesco et al. 1998).  Both AB~Aur and HD~163296 show a rich, variable, emission-line spectrum from the ultraviolet to the optical. With a few exceptions, most noticeably the observed [O\\,{\\sc i}] emission, these lines have been successfully modelled as arising in an inhomogeneous stellar wind (Catala et al. 1989; B\\\"ohm et al. 1996; Bouret et al. 1997). Infall of material has been detected in both HD~163296 and AB~Aur through monitoring of UV and optical absorption and emission lines and can be explained by the presence of infalling evaporating exocomets (Grady et al. 1996, 1999). In the infrared, both stars are among the sources with the strongest 10~$\\mu$m silicate feature in emission (Cohen 1980; Sitko 1981; Sorrell 1990). Sitko et al. (1999) compared the 10~$\\mu$m silicate feature in HD~163296 with that of solar-system comets and found a striking resemblance with that of comet Hale-Bopp. Neither star shows the 3.29~$\\mu$m UIR emission band present toward many Herbig Ae/Be stars (Brooke et al. 1993; Sitko et al. 1999). Basic astrophysical parameters of both stars are listed in Table~1.  There is strong evidence for the presence of a circumstellar disk in both AB~Aur and HD~163296. Bjorkman et al. (1995) detected a 90\\degr~flip of the polarization angle between the optical and the ultraviolet in HD~163296, which they interpreted as evidence for a flattened, disk-like structure. Mannings \\& Sargent (1997) resolved the gaseous disks surrounding AB~Aur and HD~163296 using CO millimeter wave aperture synthesis imaging. Using continuum measurements at 1.3~mm, the same authors also resolved the circumstellar dust disk of HD~163296. The AB~Aur dust disk has also been resolved in the infrared, and shows a surprisingly strong dependence of disk diameter on wavelength, ranging from 0\\farcs0065 (0.94~AU) at 2.2~$\\mu$m (Millan-Gabet et al. 1999), through 0\\farcs24 (35 AU) at 11.7~$\\mu$m to 0\\farcs49 (70~AU) at 17.9~$\\mu$m (Marsh et al. 1995).  In this paper we present new infrared spectra of AB~Aur and HD~163296 obtained with the Short- and Long Wavelength Spectrometers on board the {\\it Infrared Space Observatory} (ISO; Kessler et al. 1996). We will discuss these spectra and their implications for the evolution of dust in Herbig systems. In a subsequent paper (Bouwman et al. 2000), we will describe a model for the circumstellar dust disks of AB~Aur and HD~163296 and apply it to these data. ",
        "conclusions": "We have shown that the main difference between the AB~Aur and HD~163296 systems is that HD~163296 contains a population of very large (mm to cm-sized), cold, partially crystalline, dust grains, which is absent in AB~Aur. AB~Aur contains a population of small dust grains (PAHs), which is absent in HD~163296. In view of the fact that the stellar parameters are nearly identical except for stellar rotation, these differences are remarkable. This must mean that either the evolution of the dust composition in protoplanetary disks happens within the error in the age determination of both systems (2$^{+2}_{-1}$ Myr for AB~Aur vs. 4$^{+4}_{-2}$ Myr for HD~163296; van den Ancker et al. 1998), or that the evolution of the dust is dominated by external factors.  We have also shown that in AB~Aur and HD~163296 iron oxide is a constituent of the circumstellar dust mixture and can also be responsible for the observed excess emission near 3~$\\mu$m. Since the near-infrared excess exhibited by our programme stars is by no means unusual for a Herbig Ae/Be star, this means that the same applies for the entire group of Herbig stars. Therefore the results of models attributing this near infrared excess emission to very hot dust from an actively accreting disk (e.g. Hillenbrand et al. 1992) must be regarded with some caution. The case of HD~163296 demonstrates that infrared broad-band photometry can be completely dominated by emission from solid-state features, which must also be taken into account in any future modelling of the energy distributions of Herbig stars.  The detection of PAHs in AB~Aur shows that ground-based surveys (e.g. Brooke et al. 1993) have underestimated the fraction of Herbig stars containing PAHs. The case of HD~163296 shows that the fraction of Herbig stars showing PAH emission will not go up to 100\\%, so models depending on the presence of very small dust grains to explain the observed near-infrared excess in Herbig Ae/Be stars (Natta et al. 1993; Natta \\& Kr\\\"ugel 1995) will not be successful in all cases. We believe iron oxide to be a more plausible explanation for this near-infrared excess.  It is difficult to attribute the absence of the 3.29~$\\mu$m feature in AB~Aur to an unusual temperature of the PAHs. Since the 6.2 and 7.7~$\\mu$m C--C stretches are strong, while the bands due to C--H bonds are weak or absent, a more promising possibility seems a very low hydrogen covering factor of the PAHs in AB~Aur or the presence of a population of large ($>$ 100 C atoms) PAHs (Schutte et al. 1993). If the new 15.9~$\\mu$m UIR band in AB~Aur is also caused by PAHs, this suggests that it is also due to a C--C bond.  To gain more insight in the evolution of dust in protoplanetary disks, it is useful to compare the spectra presented here to those of other Herbig Ae stars. Except for the absence of crystalline material, the AB~Aur spectrum and energy distribution are nearly identical to those of its older counterpart HD~100546 (Malfait et al. 1998). In the case of AB~Aur (2.5~M$_\\odot$), any possible crystallization of circumstellar dust must therefore occur at a stellar age older than 2 $\\times$ 10$^6$ years. The differences between HD~100546 and HD~163296 are larger: the crystalline dust in HD~100546 is much more prominent than that in HD~163296 and the population of large cold dust grains seen in HD~163296 is absent in HD~100546.  As these cases of AB~Aur, HD~163296 and HD~100546 demonstrate, the age of the central star and the degree of crystallization do not show a one to one correspondence, but the processes of grain growth and crystallization in protoplanetary disks are also not necessarily coupled. Studying a larger sample of Herbig Ae/Be stars might shed more light on what causes this large observed diversity in dust properties in systems which appear very similar in other aspects."
    },
    "0002/astro-ph0002506_arXiv.txt": {
        "abstract": "We discuss the X-ray properties of the cooling flows in a sample of thirty highly X-ray luminous clusters of galaxies observed with the ASCA and ROSAT satellites. We demonstrate the need for multiphase models to consistently explain the spectral and imaging X-ray data for the clusters. The mass deposition rates of the cooling flows, independently determined from the ASCA spectra and ROSAT images, exhibit good agreement and exceed 1000 \\Msunpyr~in the largest systems. We confirm the presence of intrinsic X-ray absorption in the clusters using a variety of spectral models. The measured equivalent hydrogen column densities of absorbing material are sensitive to the spectral models used in the analysis, varying from average values of a few $10^{20}$ \\apc~for a simple isothermal emission model, to a few $10^{21}$ \\apc~using our preferred cooling flow models, assuming in each case that the absorber lies in a uniform foreground screen. The true intrinsic column densities are likely to be even higher if the absorbing medium is distributed throughout the clusters. We summarize the constraints on the nature of the X-ray absorber from observations in other wavebands. A significant part of the X-ray absorption may be due to dust. ",
        "introduction": "X-ray observations of clusters of galaxies show that in the central regions of most clusters the cooling time of the intracluster gas is significantly less than a Hubble time (\\eg Edge, Stewart \\& Fabian 1992; White, Jones \\& Forman 1997; Peres \\etal 1998).  This cooling is thought to lead to a slow net inflow of material towards the cluster centre; a process known as a cooling flow (see Fabian 1994 for a review). X-ray imaging data show that gas typically `cools out' throughout the central few tens to hundreds of kpc in clusters, with ${\\dot M}(r) \\approxpropto r$, where ${\\dot M}(r)$ is the integrated mass deposition rate within radius $r$ (\\eg Thomas, Fabian \\& Nulsen 1987). Spatially resolved X-ray spectroscopy has confirmed the presence of distributed cooling gas in cooling flows, with a spatial distribution and luminosity in good agreement with the predictions from the imaging data and cooling-flow models (\\eg Allen \\& Fabian 1997).  For some years, the primary uncertainty with the standard model for cooling flows was the fate of the cooled matter (\\eg see Fabian, Nulsen \\& Canizares 1991). However, the discovery of large column densities of intrinsic X-ray absorbing material associated with cooling flows observed with the Solid State Spectrometer (SSS) on the Einstein Observatory (White \\etal 1991; Johnstone \\etal 1992) opened one interesting possibility. Follow-up spatially-resolved X-ray spectroscopy with the Position Sensitive Proportional Counter (PSPC) on ROSAT (\\eg Allen \\etal 1993; Irwin \\& Sarazin 1995; Allen \\& Fabian 1997) confirmed the presence of intrinsic X-ray absorbing material in cooling-flow clusters and further showed this material to be distributed throughout, but centrally-concentrated within, the cooling flows. The X-ray data thus identify  the intrinsic X-ray absorbing material as a plausible sink for the cooled gas deposited by the cooling flows.   In this paper we examine the X-ray properties of the cooling flows in a sample of 30 of the most X-ray luminous ($L_{\\rm Bol} > 10^{45}$ \\ergps) clusters of galaxies known. Using ASCA spectra and ROSAT High Resolution Imager (HRI) data, we present independent determinations of the mass deposition rates in the cooling flows and measure the column densities of intrinsic X-ray absorbing material associated with these systems. We show that the cooling flow model provides a consistent description of the X-ray imaging and spectral data and that the observed masses of intrinsic X-ray absorbing material are in reasonable agreement with the masses expected to have been accumulated by the cooling flows over their lifetimes. We also discuss the evidence for mass deposition from the cooling flows in other wavebands.  This is the final paper in a series that has examined the impact of cooling flows on mass measurements (Allen 1998), the $kT_{\\rm X}-L_{\\rm Bol}$ relation (Allen \\& Fabian 1998a) and metallicity measurements (Allen \\& Fabian 1998b) for X-ray luminous clusters. In this paper, we describe the data reduction and analysis procedures used in these works. A large number of other, previous studies have also analysed one or more of data sets included here (see references in the papers listed above), although these studies have not, in general, examined the properties of the cooling flows in the clusters. Exceptions are the deprojection analysis of Peres \\etal (1998), and the combined ASCA/ROSAT studies of Allen \\etal (1996), Schindler \\etal (1997), B\\\"ohringer \\etal (1998) and Rizza \\etal (1998).  The cosmological parameters $H_0$=50 \\kmpspMpc, $\\Omega = 1$ and $\\Lambda = 0$ are assumed throughout. ",
        "conclusions": "The main conclusions that may be drawn from this paper may be summarized as follows: \\vskip 0.2cm  (i) We have demonstrated the need for multiphase models to consistently explain the spectral and imaging X-ray data for the CF clusters included in our study. The mass deposition rates from the cooling flows, independently inferred from multiphase analyses of the ASCA spectra and deprojection analyses of the ROSAT HRI images, exhibit good agreement, especially once the effects of intrinsic X-ray absorption have been accounted for in a consistent manner. The mass deposition rates from the largest cooling flows exceed 1000 \\Msunpyr, identifying these as some of the most massive cooling flows known.  (ii) We have confirmed the presence of intrinsic X-ray absorption in the cluster spectra using a variety of spectral models. The measured equivalent hydrogen column densities are sensitive to the spectral models used in the analysis, ranging from a few $10^{20}$ \\apc~for a simple isothermal emission model to a few $10^{21}$ \\apc~using our preferred cooling-flow models, assuming in each case that the absorber lies in a uniform foreground screen. Both the CF and NCF systems exhibit excess absorption at a similar level (on average) when analysed with the same, simple isothermal model.  (iii) The masses of X-ray absorbing material inferred to be present in the CF clusters (assuming the absorption to be due to cold gas with a covering fraction of unity) are in reasonable agreement with the masses expected to have been accumulated by the cooling flows over their lifetimes.  (iv) The ASCA spectra constrain the covering fraction of the absorbing material acting on the cooling flows to be close to (or exceed) unity. If the X-ray absorption is due to many small, similarly-sized clouds along each line of sight, intermixed with the X-ray emitting gas, then the column densities inferred from the spectral analysis may significantly underestimate the true column densities of absorbing material in the cooling flows.  (vii) We have summarized the constraints on the physical properties of the X-ray absorbing material from observations in other wavebands. Substantial 60 and 100$\\mu$m fluxes, sufficient to account for the bolometric luminosities absorbed and reprocessed within the clusters, are detected from several of the largest cooling flows. Optical observations of star formation in the cores of cooling flows also suggests that significant masses of molecular gas are present in the central few tens of kpc of many CF systems. Constraints from 21cm observations suggest that the bulk of the material accumulated within the cooling radii cannot remain in a long-lived  reservoir of atomic hydrogen. Substantial masses of molecular gas may be distributed throughout cooling flows and as yet have avoided detection, although this material must be very cold ($T \\sim 3$K), a possibility that remains controversial. The very large X-ray column densities ($\\approxgt 10^{22}$ \\apc) required if the absorbing material has a multilayer distribution are probably incompatible with a gaseous absorber. It remains possible that dust grains, either present in the material accumulated by the cooling flows or distributed more widely in a manner similar to the X-ray gas, may be responsible for much of the X-ray absorption."
    },
    "0002/astro-ph0002038_arXiv.txt": {
        "abstract": "We searched the COBE DIRBE Sky and Zodi Atlas for a wake of dust trailing Mars and for Trojan dust near Jupiter's L5 Lagrange point. We compare the DIRBE images to a model Mars wake based on the empirical model of the Earth's wake as seen by the DIRBE and place a 3-$\\sigma$ upper limit on the fractional overdensity of particles in the Mars wake of 18\\% of the fractional overdensity trailing the Earth.  We place a 3-$\\sigma$ upper limit on the effective emitting area of large (10-100 micron diameter) particles trapped at Jupiter's L5 Lagrange point of $6 \\times 10^{17}$ cm${}^2$, assuming that these large dust grains are distributed in space like the Trojan asteroids.  We would have detected the Mars wake if the surface area of dust in the wake scaled simply as the mass of the planet times the Poynting-Robertson time scale. ",
        "introduction": "We worked with a version of the DIRBE data set which contains the sky brightness with a model for the background zodiacal emission subtracted: the zodi-subtracted weekly data set from the DIRBE Sky and Zodi Atlas (DSZA).  The DIRBE team created the DZSA by fitting an 88-parameter model of the zodiacal dust emission to the observed sky brightness \\markcite{kels98}(Kelsall et al. 1998).  The model includes a smooth widened-fan component, three pairs of dust bands near the ecliptic plane, and the Earth's ring and trailing wake, but no structures associated with Mars, Jupiter, or any other planets.  The zodi-subtracted weekly data set contains 41 files for each DIRBE band, spanning a period from 10~December 1989 to 21~September 1990, covering 10 bands, centered at 1.25, 2.2, 3.5, 4.9, 12, 25, 60, 100, 140, and 240 microns.  Galactic emission dominates the zodi-subtracted maps in the mid and far-infrared near the galactic plane.  Near the ecliptic plane, the zodi-subtracted maps are dominated by residuals from the subtraction of the dust bands that are associated with prominent asteroid families \\markcite{reac97}\\markcite{kels98} (Reach et al. 1997; Kelsall et al. 1998).  The presence of these bands makes searching for smooth, faint heliocentric rings of dust near the ecliptic plane impossible.  However, we could hope to distinguish a blob of dust following a planet across the sky from other cloud components and from the galactic background by the apparent motion of the blob during the COBE mission.  The COBE satellite orbited the Earth near the day/night terminator and repeatedly mapped a swath of the sky extending about 30 degrees before and behind the terminator (see the COBE DIRBE Explanatory Supplement \\markcite{cobe}(1997) for details).  Each weekly map contains a robust average of all the week's data and covers a region a little larger than the daily viewing swath.  This weekly averaging tends to exclude transient events that would contaminate our final maps, but should not otherwise significantly affect a search for large features that move only a few degrees per week.  Figure 1 is a schematic view of the solar system during week 34 of the mission (9--16 July 1990) showing the positions of Earth, Mars and Jupiter, and the DIRBE viewing swath for that week.  Because DIRBE never imaged the sky within $60^{\\circ}$ of the sun, the orbits of Mercury and Venus, for instance, do not appear in the data. Mars appeared in the DIRBE viewing swath for 25 weeks of the mission, and moved $111^{\\circ}$ in ecliptic longitude during those weeks.  Jupiter moved only $40^{\\circ}$ in ecliptic longitude during the entire mission, but this is sufficient to allow some crude background subtraction.  More distant planets moved less.  Based on these constraints, we decided to search the weekly maps for dust features following the orbital paths of Mars and Jupiter.  Figure 2 shows the intersection of the DIRBE viewing swath with the ecliptic plane throughout the 41 weeks of the mission, and the ecliptic longitudes of Mars, Jupiter and the Sun during those weeks. ",
        "conclusions": "The zodiacal cloud near the ecliptic plane is a complex tapestry of dynamical phenomena.  We could not detect the Mars wake or Jupiter's Trojan clouds among the asteroidal dust bands in the DIRBE maps, despite the efforts of the DIRBE team to subtract these bands from the maps. We would have detected the Mars wake if it had 18\\% of the overdensity of the Earth wake, based on our empirical model for the Earth wake.  This upper limit illustrates the complexity of relating resonant structures in circumstellar dust disks to the properties of perturbing planets.  For instance, we would have detected the Mars wake if the surface area of the dust in the wake scaled simply with the mass of the planet times the Poynting-Robertson time scale.  The Trojan clouds, by our crude estimation, would have been a few orders of magnitude too faint to detect if the dust concentration in these clouds were at its mean levels.  However, a transient cloud created by a recent collision of Trojan asteroids might have been detectable.  We measured that the total 60-micron flux from large (10--100 micron diameter) dust particles trapped at Jupiter's L5 Lagrange point is less than $\\sim30$ kJy."
    },
    "0002/astro-ph0002512_arXiv.txt": {
        "abstract": "We have used the star camera on the WIRE satellite to observe the K0~III star $\\alpha$ UMa, and we report the apparent detection of 10 oscillation modes. The lowest frequency mode is at $1.82~\\mu \\rm Hz$, and appears to be the fundamental mode. The mean spacing between the mode frequencies is $2.94~\\mu \\rm Hz$, which implies that all detected modes are radial. The mode frequencies are consistent with the physical parameters of a K0~III star, if we assume that only radial modes are excited. Mode amplitudes are $100 - 400~\\mu \\rm mag$, which is consistent with the scaling relation of Kjeldsen \\& Bedding \\markcite{kandb95}(1995). ",
        "introduction": "Over the past few decades, our understanding of the interior of the Sun -- its thermodynamic structure, internal rotation, and dynamics -- has been revolutionized by the technique of helioseismology, the study of the frequencies and amplitudes of seismic waves that penetrate deep into the solar interior (Leibacher et al. \\markcite{leibacher85}1985; Duvall et al. \\markcite{duvall88}1988; Schou et al. \\markcite{schou98}1998). The high quality of the modeling in the solar case is made possible by the large number ($\\approx 10^7$) of modes visible in the Sun. Unfortunately, the lack of spatial resolution inherent in stellar observations limits the number of detectable modes in stars to only a few (those with low degree $l$). Nonetheless, successful detection of even a few modes has the potential to provide greatly improved values for fundamental stellar parameters such as mass, abundance, and age (Gough \\markcite{gough87}1987; Brown et al. \\markcite{brown94}1994)  While oscillations have been successfully detected on roAp stars, $\\delta$ Scuti stars, and white dwarfs, these stars all show oscillation amplitudes several orders of magnitude larger than is expected from solar-like stars. More recently, however, a number of authors (Hatzes \\& Cochran \\markcite{hatzes94a}1994a, \\markcite{hatzes94b}1994b; Edmonds \\& Gilliland \\markcite{edmonds96}1996) have reported the detection of periodic variability at levels of a few hundred $\\rm m~s^{-1}$ and/or several millimagnitudes in a number of K giants. However, each detection is of only a single mode, and multimodal oscillations have yet to be unambiguously detected in any cool star other than the Sun.  In March 1999, NASA launched the Wide-Field Infrared Explorer satellite, with the intent of carrying out an infrared sky survey to better understand galaxy evolution. Unfortunately, within days of launch the primary science instrument on WIRE failed due to loss of coolant. However, the satellite itself continues to function nearly perfectly, and in May we began a program of asteroseismology using the WIRE's onboard 52 mm aperture star camera. Below we report WIRE's probable detection of multimodal oscillations on a cool star, which is the first of its kind. ",
        "conclusions": ""
    },
    "0002/astro-ph0002348_arXiv.txt": {
        "abstract": "The X-ray source 4U1700+24/HD154791 is one of the few galactic sources whose counterpart is an evolved M star \\cite{ddf,garcia,gaudenzi}. In X-rays the source shows extreme erratic variability and a complex and variable spectrum. While this strongly suggests accretion onto a compact object, no clear diagnosis of binarity was done up to now. We report on \\A and \\B X--ray broad band observations of this source and on ground optical observations from the Loiano 1.5 m telescope. ",
        "introduction": "In optical astronomy the identification of a binary system comes in most cases from the observation of photometric and/or radial velocity variations. As not all X-ray binaries have known optical counterparts, a further effective criterium in galactic X--ray astronomy for the identification of a binary system with an accreting compact object was often based on the observed X--ray luminosity. For X--ray binaries harbouring a neutron star or possibly a black hole, luminosities \\lx of the order of 10$^{34}$ -- 10$^{35}$ \\ergs are easily reached. The diagnosis of the presence of a neutron star in most cases is directly confirmed by the observation of pulsations or thermonuclear bursts, apart from bright persistent Low Mass X--Ray Binaries (LMXRBs). X--ray binaries harbouring white dwarfs also show some distinctive features. As an example in polars and intermediate polars optical and UV observations often reveal the distinctive signatures of the presence of a white dwarf in the system. Orbital periods and light curves also add unambiguous and reliable evidence of the presence of white dwarfs in this class of X--ray binaries.\\\\ For a number of X--ray sources the identification of a class or even the diagnosis of binarity is rather difficult, especially when the observed X--ray luminosity is $\\leq 10^{33}$ \\ergs.\\\\ \\src belongs to this class. The optical counterpart was identified by Garcia et al. \\cite{garcia} as a late type giant on the basis of the positional coincidence with a HEAO1--A3 error box. The optical spectrum of this giant looks quite normal \\cite{ddf,garcia}, even if Gaudenzi and Polcaro \\cite{gaudenzi} find some interesting and variable features in its spectrum. Variable UV line emission was detected \\cite{ddf,garcia} in different IUE pointings, showing at last some unusual features in the emission from this otherwise normal giant. These high excitation lines are likely linked to the same mechanism that produces the observed X--ray emission. In spite of various attempts, no evidence of a binary orbit was obtained from radial velocity analysis of optical spectra.\\\\ The X--ray source shows extreme erratic variability, but no pulsations were detected. The rapid (10--1000 s) time variability is strongly suggestive of turbulent accretion, often observed in X--ray binaries. The X--ray spectrum is rather energetic and was measured up to 10 keV. The X--ray luminosity \\lx $\\sim 10^{33}$ \\ergs at an assumed distance of 730 pc \\cite{garcia} may be marginally consistent with coronal emission, even if an evolved giant is not expected to be a strong X-ray emitter. Therefore the picture emerging from observations gives only hints in favour of a binary system, given that no ``classical'' feature to be associated to the presence of a compact object was found.\\\\ We have observed this source for $\\sim$15 years both with X--ray satellites (EXOSAT, \\A and \\B) and with ground optical observations from the Loiano 1.5 m and 0.6 m telescopes of the Bologna Astronomical Observatory. Here we report on the \\A and \\B observations, performed respectively on March 8, 1995 and on  March 27, 1998. We also report on photometric optical UBVRI monitoring. ",
        "conclusions": "The observations we report still cannot be used to perform a ``classical'' diagnosis of binarity. We nevertheless note some interesting similarities with other low luminosity X--ray sources. In particular some interesting similarities can be found with the X--ray emission from $\\gamma$ Cas. The power spectrum is strikingly similar and the energy spectrum shows a similar shape, even if no iron line is detected in \\ssrc.  However this close resemblance of the properties of the X--ray emission does not help to determine the presence of a compact object in a binary system, as for $\\gamma$ Cas itself the diagnosis of binarity is not completely assessed. In fact Owens \\etal \\cite{owens} favour the hypothesis of a WD binary, but a completely different point of view is based on recent UV/X--ray observations of $\\gamma$ Cas \\cite{smith}. Smith \\etal support the hypothesis that the X--ray emission of $\\gamma$ Cas comes from continuous flaring from the Be star. This hypothesis cannot be easily adapted to the case of \\src, as the much colder M giant star should not be expected to have strong and persistent X--ray flaring activity. If this is the case, i.e. the observed X--ray emission from \\ssrc is coronal, HD154791 should be an exception in its own class. If the similarity of the properties of the X--ray emission from \\ssrc and $\\gamma$ Cas comes from a common origin, we suggest that the WD binary hypothesis is much more comfortable and more easily met.  {\\bf Acknowledgements}. This research is supported by the Agenzia Spaziale Italiana (ASI) and the Consiglio Nazionale delle Ricerche (CNR) of Italy. \\B is a joint program of ASI and of the Netherlands Agency for Aerospace Programs (NIVR). The \\A observation was performed as part of the joint ESA/Japan scientific program. CB, AG and AP acknowledge a grant from ``Progetti di ricerca ex-quota 60\\%'' of Bologna University."
    },
    "0002/astro-ph0002381_arXiv.txt": {
        "abstract": " ",
        "introduction": "Was the primordial universe turbulent or non-turbulent soon after the Big Bang? How  did the hydrodynamic state of the early universe affect the formation of structure from gravitational forces, and how did the formation of structure by gravity affect the hydrodynamic state of the flow?  What can be said about the dark matter that comprises $99.9 \\%$ of the mass of the universe according to most cosmological models?  Space telescope measurements show answers to these questions persist literally frozen as fossils of the primordial turbulence and nonturbulence that controlled structure formation, contrary to standard cosmology which relies on the erroneous Jeans 1902 linear-inviscid-acoustic theory and a variety of associated misconceptions (e. g., cold dark matter). When effects of viscosity, turbulence, and diffusion are included,  vastly different structure scenarios and  a clear explanation for the dark matter emerge~\\cite{gib96}.  From Gibson's 1996 theory the baryonic (ordinary) dark matter is comprised of proto-globular-star-cluster (PGC) clumps of hydrogenous planetoids termed ``primordial fog particles''(PFPs), observed by Schild 1996 as ``rogue planets ... likely to be the missing mass'' of a quasar lensing galaxy~\\cite{sch96}.  The weakly collisional non-baryonic dark matter diffuses to form outer halos of galaxies and galaxy clusters~\\cite{tys95}. \\\\ ",
        "conclusions": "The Figure summarizes the evolution of structure and turbulence in the early universe, as inferred from the present  nonlinear fluid mechanical theory. It is very different, very early, and very gentle compared to the standard model, where structure formation in baryonic matter is forbidden in the plasma epoch because $L_J$ is larger than $L_H = ct$ and galaxies collapse at 140 million years (redshift z=20) producing $10^{36}$ kg Population III superstars that explode and re-ionize the universe to explain the missing gas (sequestered in PFPs).  No such stars, no galaxy collapse, and no re-ionization occurs in the present theory.  To produce the structure observed today, the concept ``cold dark matter'' (CDM) was invented; that is, a hypothetical non-baryonic fluid of ``cold'' (low speed) collisionless particles with adjustable $L_J$ small enough to produce gravitational potential wells to drive galaxy collapse.   Cold dark matter is unnecessary in the present theory. Even if it exists it would not behave as required by the standard model.  Its necessarily small collision cross section requires $L_{SD} \\gg L_J$ so it would diffuse out of its own well, without fragmentation if $L_{SD} \\gg L_H$. The immediate formation of ``primordial fog particles'' from all the neutral gas of the universe emerging from the plasma epoch permits their gradual accretion to form the observed small ancient stars in dense globular-star-clusters known to be only slightly younger than the universe.  These could never form in the intense turbulence of galaxy collapse in the standard model because $L_{ST}$ scales would be too large. \\\\  \\begin{figure}[!h] \\vspace{1mm} \\begin{center} \\includegraphics[width= 0.8 \\linewidth]{etc8_fig.eps} \\end{center}  \\vskip-2.8in  \\caption{Evolution of structure and turbulence in the early universe} \\end{figure}"
    },
    "0002/astro-ph0002454_arXiv.txt": {
        "abstract": " ",
        "introduction": "Previously (Maller et~al. 1999, hereafter MSPP) we have shown that it is possible to account for the kinematic properties of damped Lyman alpha systems (DLAS) as measured by Prochaska \\& Wolfe (1997, 1998, hereafter PW97 and PW98) in the context of semi-analytic models (SAMs) (Somerville \\& Primack 1999, hereafter SP99). In these models, hierarchical structure formation is approximated by constructing a merger tree for each dark matter halo. A natural consequence is that every virialized halo may contain not only a central galaxy, but also a number of satellite galaxies as determined by its merging history.  Thus the kinematics of the DLAS arise from the combined effects of the internal rotation of gas disks and the motions between gas disks within a common halo.  Here we investigate the sensitivity of this model to some of the assumptions made in MSPP, including the modeling of satellite dynamics, the scale height of the gas, and the cosmology. ",
        "conclusions": "In MSPP we argued that the observed kinematics of DLAS can be reproduced in hierarchical models if a significant fraction of the lines of sight pass through multiple disks orbiting within a common dark matter halo. However, we found that in order to obtain a high enough cross section for multiple hits, we had to assume very large radial extents for the gaseous disks in our models.  We have tested the robustness of these conclusions by modifying several of the uncertain ingredients of our models. We find that modifying the dynamical friction timescale of the satellites by assuming a different intial radius or orbit has a small effect on our results. Similarly, our results do not seem to be sensitive to the assumed cosmology. Increasing the vertical scale height of the disks has a larger effect and leads to models that are more physically plausible and still produce good agreement with the diagnostic statistics of PW. Our results suggest that in order to reconcile the observed kinematics of DLAS with hierarchical theories of structure formation, gaseous disks at high redshift must be large in radial extent and thickened or warped. The physical cause of these properties remains obscure, however we speculate that tidal encounters or outflows could be responsible."
    },
    "0002/astro-ph0002104_arXiv.txt": {
        "abstract": " ",
        "introduction": "The Current System}  Currently the ADS system consists of four semi-autonomous (to the user) abstract services covering Astronomy, Instrumentation, Physics, and Astronomy Preprints.  Combined there are nearly 1.5 million abstracts and bibliographic references in the system.  The Astronomy Service is by far the most advanced, and accounts for $\\sim 85$\\%\\ of all ADS use; it ought be noted, however, that the Instrumentation Service contains more abstracts than Astronomy, and a subset of that service is used by the Society of Photo-Optical Industrial Engineers as the basis of the official SPIE publications web site.  All of what follows will refer only to the Astronomy service.  \\subsection{Data}  Here is a brief overview of the data in the ADS system, a complete description is in DATA.  \\subsubsection{Abstracts}  The ADS began with the abstracts from the NASA STI database, in printed form these abstracts were the union of the International Aerospace Abstracts and the NASA Scientific and Technical Abstracts and Reports (NASA STAR).  While the STI branch has had to substantially cut back on their abstracting of the journal literature, we still get abstracts of NASA reports and other materials from them.  We now receive basic bibliographic information (title, author, page number) from essentially every journal of astronomy.  Most also send us abstracts, and some cannot send abstracts, but allow us to scan their journals, and we build abstracts through optical character recognition.  Finally we receive some abstracts from the editors of conference proceedings, and from individual authors.  The are $\\sim$500,000 different astronomy articles indexed in the ADS, the database is nearly complete for the major journals articles beginning in 1975.  \\subsubsection{Bitmaps}  The ADS has obtained permission to scan, and make freely available on-line, page images of the back issues of nearly all of the major journals of astronomy.  In most cases the bitmaps of current articles are put on-line after a waiting period, to protect the financial integrity of the journal.  DATA describes the current status of these efforts.  We plan to provide for each collaborating journal, in perpetuity, a database of page images (bitmaps) from volume 1 page 1 to the first issue which the journal considers to be fully on-line as published. This will provide (along with the indexing and the more recent archives held by the journals) a complete electronic digital library of the major literature in astronomy.  On a longer term we plan to scan old observatory reports, and defunct journals, to finally have a full historical collection on-line.  This work is beginning with a collaboration with the Harvard Preservation Project (\\cite{1997AAS...191.3502E}; \\cite{1995VA.....39..161C}).  \\subsubsection{Links}  ADS responds to a query with a list of references and a set of hyperlinks showing what data is available for each reference.  There are $\\sim$1.73 million hyperlinks in the ADS, of which $\\sim$ 31\\%\\ are to sources external to the ADS project.  The largest number of external links are to SIMBAD, NED, and the electronic journals.  A rapidly growing number, although still small in comparison to the others, are to data tables created by the journals and maintained by the CDS and the ADC at Goddard.  SEARCH describes the system of hyperlinks in detail.  \\subsubsection{Citations and References}  The use of citation histories is a well known and effective tool for academic research (\\cite{1979cita.book.....G}); their inclusion in the ADS has been planned since the conception of the service.  In 1996 the AAS purchased a subset of the Science Citation Index from the Institute for Scientific Information, to be used in the ADS; this was updated in 1998.  This subset only contains references which were already in the ADS, thus it is seriously incomplete in referring to articles in the non-astronomical literature.  This citation information currently spans January 1982-September 1998.  The electronic journals all have machine readable, web accessible, reference pages.  The ADS points to these with a hyperlink where possible.  Several publishers allow us to use these to maintain citation histories; we do this using our reference resolver software (see ARCHITECTURE).  The same software is also used by some publishers to check the validity of their references, pre-publication.  Additionally we use optical character recognition to create reference and citation lists for the historical literature, after it is scanned (\\cite{1999AAS...195.8209D}).  \\subsubsection{Collaboration with CDS/SIMBAD}  The Strasbourg Data Center (CDS) has long maintained several of the most important data services for astronomy (e.g. \\cite{1971BICDS...1....2J}; \\cite{1973BICDS...4...27J}; \\cite{1998adass...7..470G}); access to parts of the CDS data via ADS is a key feature of the ADS.  ADS users are able to make joint queries of the ADS bibliographic database and the CDS/SIMBAD bibliographic data base.  When SIMBAD contains information on a object which is referred to in a paper whose reference is returned by ADS then ADS also returns a pointer to the SIMBAD data.  When a paper has a data table which is kept on-line at the CDS the ADS returns a pointer to it.  The CDS-ADS collaboration is at the heart of Urania (section \\ref{Urania}).  More recently ADS has entered into a collaboration with the National Extragalactic Database (NED; \\cite{1988alds.proc..335H}, \\cite{1992adass...1...47M}) which is similar to the SIMBAD portion of the CDS-ADS collaboration.  \\subsection {Search Engine}  The basic design assumption behind the search engine, and other user interfaces, is that the user is an expert astronomer.  This differs from the majority of information retrieval systems, which assume that the user is a librarian.  The default behavior of the system is to return more relevant information, rather than just the most relevant information, assuming that the user can easily separate the wheat from the chaff.  In the language of information retrieval this is favoring recall over precision.  SEARCH describes the user interface in detail.  \\subsection{Hardware and Software Architecture}  The goals of our hardware and software systems are speed of information delivery to the user, and ease of maintainability for the staff.  We thus pre-compute many things during our long indexing process for later use by the search engine; we have highly optimized all code which is run by user processes; we have developed a worldwide network of mirror sites to speed up internet access.  ARCHITECTURE describes these systems.  \\subsection {Data Ingest}  The basic rule for what books and periodicals the ADS covers is: if it is in the Center for Astrophysics library it should be in the ADS.  As a goal we are still some ways from realization.  We have recently adopted a second rule for inclusion: if it is referenced by an article in a major scholarly journal of astronomy it should be in the ADS.  DATA describes the ADS coverage, and ingest procedures. ",
        "conclusions": ""
    },
    "0002/astro-ph0002042_arXiv.txt": {
        "abstract": "We investigate the origin of the variation of the gas mass fraction in the core of galaxy clusters, which was indicated by our work on the X-ray fundamental plane. Applying a spherical collapse model of cluster formation and considering the effect of shocks on preheated intracluster gas, we construct a simple model to predict the spatial gas distribution of clusters. As is suggested by our previous work, we assume that the core structure of clusters determined at the cluster collapse has not been much changed after that. The adopted model supposes that the gas distribution characterized by the slope parameter is related to the preheated temperature. Comparison with observations of relatively hot ($\\gtrsim 3$ keV) and low redshift clusters suggests that the preheated temperature is about 0.5-2 keV, which is higher than expected from the conventional galactic wind model and possibly suggests the need for additional heating such as quasars or gravitational heating on the largest scales at high redshift. The dispersion of the preheated temperature may be attributed to the gravitational heating in subclusters. We calculate the central gas fraction of a cluster from the gas distribution, assuming that the global gas mass fraction is constant within a virial radius at the time of the cluster collapse. We find that the central gas density thus calculated is in good agreement with the observed one, which suggests that the variation of gas mass fraction in cluster cores appears to be explained by breaking the self-similarity in clusters due to preheated gas. We also find that this model does not change major conclusions on the fundamental plane and its cosmological implications obtained in previous papers, which strongly suggests that not only for the dark halo but also for the intracluster gas the core structure preserves information about the cluster formation. ",
        "introduction": "Correlations among physical quantities of clusters of galaxies are very useful tools for studying the formation of clusters and cosmological parameters. Recently, we have found that clusters at low redshifts ($z\\lesssim 0.1$) form a plane (the fundamental plane) in the three dimensional space represented by their core structures, that is, the central gas density $\\rho_{\\rm gas,0}$, core radius $r_c$, and X-ray temperature $T_{\\rm gas}$ \\citep[Paper~I]{fuj99a}. On the other hand, a simple theoretical model of cluster formation predicts that clusters should be characterized by the virial density $\\rho_{\\rm vir}$ (or the collapse redshift $z_{\\rm coll}$) and the virial mass $M_{\\rm vir}$ (\\citealt{fuj99b}, Paper~II). Thus, assuming the similarity of the dark matter distributions, clusters should form a plane in the three dimensional space of the dark matter density in the core $\\rho_{\\rm DM, c}$, the core radius of dark matter distribution $r_{\\rm DM, c}$, and the virial temperature $T_{\\rm vir}$ \\footnote{In Paper~I, we used the terms ' virial density', 'virial radius', and ' virial mass' to denote the dark matter density in the core, the core radius of dark matter distribution, and the core mass, respectively. This is because we assumed that the dark matter density in the core is proportional to the average dark matter density over the whole cluster (Paper II). To avoid possible confusions, here we use the term ' dark matter', and the term ' virial' will be used to represent spatially averaged quantities of gravitational matter (mostly dark matter) within the virialized region.}. However, the relations between the two planes are not simple; for example, it is found that $\\rho_{\\rm gas, 0}$ is not proportional to $\\rho_{\\rm DM, c}$. In Paper~I, we found that the ratio $\\rho_{\\rm gas,0}/\\rho_{\\rm DM, c}$ is not constant but obeys the relation of $\\rho_{\\rm gas,0}/\\rho_{\\rm DM, c} \\propto \\rho_{\\rm DM, c}^{-0.1}M_{\\rm DM, c}^{0.4}$, where $M_{\\rm DM, c}$ is the core mass. This raises the question how the segregation between gas and dark matter occurs.  In the hierarchical structure formation, dark halos are expected to obey scaling relations. In fact, numerical simulations suggest that the density distribution in dark halos take a universal form as claimed by \\citet{nav96, nav97}. On a cluster scale, it can be approximated by $\\rho_{\\rm DM}(r)\\propto r^{-2}$ for $r\\lesssim 1$ Mpc, where detailed X-ray observations have been done \\citep{mak98}. On the contrary, observations show that the slope of the density profile of the hot diffuse intracluster gas has a range of value. Radial surface brightness profiles of X-ray emission are often fitted with the conventional $\\beta$ model as \\begin{equation} \\label{eq:sur} I(R)=\\frac{I_0}{(1+R^2/r_c^2)^{3\\beta_{\\rm obs}-1/2}}\\;, \\end{equation} where $\\beta_{\\rm obs}$ is the slope parameter \\citep{cav78}. If the intracluster gas is isothermal, equation (\\ref{eq:sur}) corresponds to the gas density profile of \\begin{equation} \\label{eq:gas_obs} \\rho_{\\rm gas}(r)=\\frac{\\rho_{\\rm gas, 0}}{(1+r^2/r_c^2) ^{3\\beta_{\\rm obs}/2}}\\:. \\end{equation} Observations show that the slope parameter takes a range $\\beta_{\\rm obs} \\sim 0.4-1$ \\citep{jon84, jon99}. This means that for $r>>r_{\\rm c}$, the density profiles range from $\\propto r^{-1.2}$ to $\\propto r^{-3}$, which are more diverse than those of dark matter. Moreover, observations show that the clusters with large $r_c$ and $T_{\\rm gas}$ tend to have large $\\beta_{\\rm obs}$ (e.g. \\citealt{neu99}; \\citealt{hor99}; \\citealt{jon99}). Since the average gas fraction of clusters within radii much larger than $r_{\\rm c}$ should be universal and the dark matter distribution of clusters is also universal, the variation of $\\beta_{\\rm obs}$ is expected to correlate with that of the gas fraction in the core region. In other words, the gas fraction at the core is not the same as that of the whole cluster and is not proportional to the dark matter density. This fact must be taken care of when we discuss cosmological parameters using observational X-ray data. Since the emissivity of X-ray gas is proportional to $\\rho_{\\rm gas}^2$, most of the X-ray emission of a cluster comes form the central region where $\\rho_{\\rm gas}$ is large. Although in Papers I and II, we did not take account of the effects of $\\beta_{\\rm obs}$, we did find the gas mass fraction in the core region is diverse by analyzing the X-ray emission. In this paper, we reanalyze the data taking account of $\\beta_{\\rm obs}$ and discuss the relation between core and global gas mass fractions. We will also show that major conclusions on the fundamental relations are not changed.  The variation of gas mass fraction itself has been investigated by several authors (e.g. \\citealp{ett99, arn99}). \\citet{ett99} argue that it is partially explained if the dark matter has a significant baryonic component. Another possible explanation of the diverse gas distributions is that intracluster gas had already been heated before the collapse into the cluster; the energetic winds generated by supernovae are one possible mechanism to increase gas entropy (e.g. \\citealp{dek86, mih94}). In fact, \\citet{pon99} find that the entropy of the intracluster gas near the center of clusters is higher than can be explained by gravitational collapse alone. In order to estimate the effect of the preheating on intracluster gas, we must take account of shocks forming when the gas collapses into the cluster; they supply additional entropy to the gas. \\citet{cav97, cav98, cav99} have investigated both the effects and predicted the relation between X-ray luminosities and temperatures ($L_{\\rm X}-T_{\\rm gas}$ relation). They predicted that the gas distributions of poor clusters are flatter than those of rich clusters, which results in a steeper slope of $L_{\\rm X}-T_{\\rm gas}$ relation for poor clusters. This is generally consistent with the observations. It is an interesting issue to investigate whether this scenario provides a natural explanation for the observed dispersion of gas mass fraction in the cluster core and whether it reproduces the X-ray fundamental plane we found in Paper I in our general theoretical scenario.  In order to clarify what determines the gas distribution, we construct as a simple model as possible. Although many authors have studied the preheating of clusters (\\citealt{kai91,evr91,met94,bal99,kay99,wu99,val99}), this is the first time to consider the influence of the preheating and shocks on the fundamental plane and two-parameter family nature of clusters paying attention to the difference between the collapse redshift $z_{\\rm coll}$ and the observed redshift $z_{\\rm obs}$ of clusters explicitly. In \\S\\ref{sec:model}, we explain the model of dark matter potential and shock heating of intracluster gas. In \\S\\ref{sec:result}, we use the model to predict $\\beta_{\\rm obs}-T_{\\rm gas}$ and $\\beta_{\\rm obs}-r_{\\rm c}$ relations, and the fundamental plane and band of clusters. The predictions are compared with observations. ",
        "conclusions": "We have investigated the influence of heating before cluster collapse and shocks during cluster formation on the gas distribution in the central region of clusters of galaxies. We assumed that the core structure has not much changed since the formation of a cluster. Using a spherical collapse model of a dark halo and a simple shock model, we predict the relations among the slope of gas distribution $\\beta$, the gas temperature $T_{\\rm gas}$, and the core radius $r_{\\rm c}$ of clusters. By comparing them with observations of relatively hot ($\\gtrsim 3$ keV) and low redshift clusters, we find that the temperature of the preheated gas collapsed into the clusters is about $0.5-2$ keV. Since the temperature is higher than that predicted by a preheating model of supernovae, it may reflect the heating by quasars or gravitational heating on the largest scales at high redshift. Moreover, gravitational heating in subclusters assembled when the clusters formed also seems to affect the temperature of the preheated gas and produce the dispersion in the preheating temperature. Assuming that the global gas mass fraction of clusters are constant, we predict that the gas mass fraction in the core region of clusters should vary correlating with $\\beta$ through a simple law, which is shown to be consistent with the observations. Thus, we conclude that the variation of the gas mass fraction in the cluster core is due to the shock heating of preheated gas. Furthermore, we have confirmed that the observed fundamental plane and band of clusters are reproduced by the model even when the effects of preheating are taken into account. Thus, major conclusions about the cluster formation and cosmology obtained in our previous papers are not changed."
    },
    "0002/astro-ph0002364_arXiv.txt": {
        "abstract": "Cosmological linear perturbation theory predicts that the peculiar velocity $\\veV(\\vex)$ and the matter overdensity $\\delta(\\vex)$ at a same point $\\vex$ are statistically independent quantities, as log as the initial density fluctuations are random Gaussian distributed. However nonlinear gravitational effects might change the situation. Using  framework of second-order perturbation theory and the Edgeworth expansion method, we study local density dependence of bulk velocity dispersion that is coarse-grained at a weakly nonlinear scale. For a typical CDM model, the first nonlinear correction of this constrained bulk velocity dispersion amounts to  $\\sim 0.3\\delta$ (Gaussian smoothing)  at a  weakly nonlinear scale with a very  weak dependence on cosmological parameters.     We also compare our analytical prediction with published numerical results given at nonlinear regimes. ",
        "introduction": "The peculiar velocity field is one of the most fundamental quantities to analyze the large-scale structure in the universe ({\\it e.g.} Peebles 1980). It is considered to  reflect dynamical nature of density fluctuations of gravitational matter. The peculiar velocity field is usually observed  using astrophysical objects ({\\it e.g.} galaxies), as  determination of distances is  crucial for measuring  peculiar velocities (Dekel 1994, Strauss \\& Willick 1995). There is a possibility that statistical aspects of the  velocity field traced by these objects and that traced  by dark matter particles might be different.  This difference is generally called ``velocity bias'' and its elucidation becomes highly important in observational cosmology ({\\it e.g.} Cen \\& Ostriker 1992, Narayanan, Berlin \\& Weinberg 1998, Kaufmann et al. 1999).  Velocity bias is often discussed numerically with making ``galaxy particles\" in some effective manners. But here we discussed  a more basic phenomenon. It is known that statistics of the peculiar velocity field depend largely on local density contrast. For example, both the single particle and pairwise velocity dispersions   of dark matter particles are known to be increasing function of local density (Kepner, Summers, \\& Strauss 1997, Strauss, Cen \\& Ostriker 1998, Narayanan et al. 1998). Analysis of pairwise velocity statistics  is  interesting from theoretical point of views, and also very important in observational cosmology (Peebles 1976, Davis \\& Peebles 1983, Zurek et al. 1994, Fisher et al. 1994,  Sheth 1996, Diaferio \\& Geller 1996, Suto \\& Jing 1997, Seto \\& Yokoyama 1998a, 1999, Jing \\& B$\\ddot{\\rm o}$rner 1998, Juszkiewicz, Fisher \\& Szapudi 1998, Seto 1999b,  Juszkiewicz, Springel \\& Durrer 1999). But we do not discuss it here and concentrate on velocity field characterized by single point  which is simpler to analyze theoretically.  Linear perturbation theory predicts that the peculiar velocity $\\veV(\\vex)$ and the density contrast $\\delta(\\vex)$ at a given point $\\vex$ is statistically independent,  as long as the  initial density fluctuations are  random Gaussian distributed. Namely, the joint probability distribution function $P(\\veV, \\delta)$ can be written in a  form as $P_1(\\veV)P_2(\\delta)$.  It is not surprising that the  peculiar velocity of each particle is largely affected by nonlinear gravitational effects and shows local density dependence described above.  But what can we expect for the smoothed (bulk) velocity that is field coarse grained at  some spatial scale $R$?  Due to nonlinear mode couplings, the relation   $P_1(\\veV)P_2(\\delta)$ valid for linear theory must be modified and  bulk velocity dispersion must also  depend on local density contrast defined at the same  smoothing scale $R$ (see Bernardeau 1992, Chodorowski \\& {\\L}okas 1997, Bernardeau et al. 1999 for the velocity divergence field).  However, Kepner, Summers \\& Strauss (1997) showed from cold-dark-matter (CDM) and hot-dark-matter (HDM) N-body simulations  that at nonlinear scales  ($0.77\\hmpc\\le R \\le 4.88\\hmpc$), such a local density dependence was not observed (see Figs.2(a) and 3(a) of their paper). This is an interesting contrast to the behavior of velocity field traced by each particle, as described before (Kepner et al. 1997, Narayanan et al. 1998).   In this article, we investigate local density dependence of smoothed (bulk) velocity dispersion using framework of second-order perturbation theory. We calculate the first-order nonlinear correction of the constrained velocity dispersion.  Our target is weakly nonlinear scale and somewhat larger than scale analyzed by Kepner et al. (1997). Since current  survey depth of the cosmic  velocity field is highly limited, our constrained statistics might not be useful in observational cosmology at present ({\\it e.g.} Seto \\& Yokoyama 1998b, see also Seto 1999a). Our interest in this article  is theoretically motivated one about  nonlinear gravitational dynamics.   As  the peculiar velocity field is more weighted to large-scale fluctuations (smaller wave number $k$) than the density field, perturbative treatment of smoothed velocity field would  be reasonable at weakly nonlinear scale.  Actually, Bahcall, Gramann \\& Cen (1994) showed that smoothed unconstrained velocity dispersion  in N-body simulations are well predicted by linear theory even at smoothing scale $R=3\\hmpc$ (see their Table 1).  Second-order analysis by Makino, Sasaki \\& Suto (1992) also gives consistent results to their simulations. ",
        "conclusions": "It is commonly accepted that the large-scale structure observed today is formed by gravitational instability from small primordial density fluctuations (Peebles 1980). In this picture, the peculiar velocity  and the density contrast  are fundamental quantities to characterize inhomogeneities in the universe. Linear analysis  of  cosmological perturbation theory predicts that, as long as initial fluctuations are random Gaussian distributed,  the one-point PDF of the velocity field $\\veV(\\vex)$ is statistically independent of the local density contrast   $\\delta(\\vex)$. This is an important aspect of cosmological    perturbation.  However nonlinear gravitational evolution changes the situation. Due to nonlinear mode-couplings, the  peculiar velocity field is no longer statistically independent of the  local density field. Here we have investigated bulk velocity dispersion ($\\Sigma_V^2(\\delta)$) as a function of the local density contrast  and calculated its first nonlinear correction using framework of second-order perturbation theory.  Our target has been set at weakly nonlinear regimes  where perturbative treatment must be reasonable. At present,  survey depth of velocity field is highly limited and our constrained statistics might  not be directly useful for observational cosmology. However, we believe that our theoretical study is important to understand one interesting aspect of the cosmic velocity field peculiar to its nonlinear evolution.  We have shown  that the first nonlinear correction of $\\Sigma_V^2(\\delta)$ is proportional to the local density and strongly depends on the matter power spectra, but weakly on the cosmological parameters $\\Omega$ and $\\lambda$. For typical CDM model with primordially scale-invariant fluctuations, this first-order correction is about $0.3\\delta$.  If we use  overdensity of galaxies $\\delta_g$, this correction term becomes $0.3 \\delta_g/b_1$ (the factor $b_1$ is defined in eq.[35]). This  dependence might be used to constrain the linear bias parameter $b_1$ itself (not in the usual form $\\Omega^{0.6}/b_1$) in future peculiar velocity surveys. We have also shown that the constrained  one-point PDF of velocity field keeps the Gaussian shape up to  second-order of perturbation. First nonlinear effects are completely absorbed to the velocity dispersion $\\Sigma_V^2(\\delta)$.  Numerical results by Kepner et al. (1997) have been compared with our analytical results. Their results  show almost no $\\delta$-dependence, contrary to ours. However spatial scale of their analysis (strongly nonlinear regime) is largely different from ours (weakly nonlinear regime). In the forthcoming paper, detailed numerical analysis will be presented for various smoothing lengths $R$ and power spectra (see also Seto 2000). Validity of second-order  analysis as well as  the Edgeworth expansion method for velocity vector would be also investigated numerically."
    },
    "0002/astro-ph0002478_arXiv.txt": {
        "abstract": "We present ISO observations of the S(0), S(1), S(2), and S(3) rotational lines of molecular hydrogen from  two molecular clouds near the Galactic Center (GC). We have also measured continuum dust emission at infrared wavelengths with ISO and the rotational radio lines J=1--0 of $^{13}$CO and C$^{18}$O and J=2--1 of  C$^{18}$O with the IRAM-30m telescope.  Using the dust continuum spectra and the CO lines we derive a total visual extinction of $\\sim$ 15-20 magnitudes toward these GC clouds. After correcting the H$_2$ data for extinction, the gas temperatures are $\\sim$ 250~K and the column densities of warm gas are $\\sim 2\\times 10^{21}$~cm$^{-2}$. This is the first direct measure of the H$_2$ column densities of the warm component; with this, we  estimate an NH$_3$ abundance in the warm gas of $\\sim 2~10^{-7}$. The column density of warm gas is, at least, a factor of 100 larger than the corresponding column densities derived from the  warm dust. The observed ortho-to-para ratio (OTPR) is  $\\sim$ 1, clearly below the local thermodynamical equilibrium (LTE) OTPR for gas at 250 K of $\\sim 3$. Low velocity shocks ($\\sim$ 10 km s$^{-1}$) are the most likely explanation for the column densities of warm gas and dust and the non-LTE H$_2$ OTPR. ",
        "introduction": "The central $\\sim 6^{\\circ}$ of our galaxy exhibit a large accumulation of molecular material which is forming big molecular clouds whose masses and sizes are so large as 10$^6$ M$_{\\odot}$ and 15~pc, respectively. These clouds are denser (average densities of 10$^4$ \\cmmt), more turbulent (line widths of $\\sim$ 20 \\kms), and hotter (with a warm component with temperatures, $T$, up to 200-300 K) than the clouds of the disk of the galaxy (see e.g. Morris \\& Serabyn \\cite{morris}). The high temperatures in Galactic Center (GC) clouds were known basically by observations of \\NH3 inversion lines over limited regions (G\\\"usten et al. \\cite{gusten81}; Mauersberger et al. \\cite{mauersberger}). H\\\"uttemeister et al. (\\cite{huttemeister93}) analyzed  36 molecular clouds distributed all along the Central Molecular Zone and the ``Clump 2\" complex; they showed that high kinetic temperatures  are a general characteristic of the GC clouds and not only of those located close to Sgr A and Sgr B2. In the disk of the galaxy, kinetic temperatures higher than 100 K are associated with infrared sources, that is, embedded stars which heat the dust and subsequently the gas by collisions with the dust grains. The typical sizes of such  regions are less than 1 pc. The high  kinetic temperatures in the GC clouds are found in      regions of $\\sim$ 10 pc, where one measures large column densities of cold dust ($T <$ 30 K, Odenwald \\& Fazio \\cite{odenwald}; Cox \\& Laureijs \\cite{cox}). This rules out gas-dust collisions as a possible heating mechanism of the warm component. Dissipation of turbulence due to shocks induced by the rotation of the galaxy  could be the main heating mechanism in the GC clouds (Wilson et al. \\cite{wilson82}).  Unfortunately the NH$_3$ abundance  in the warm component  was  unknown since  one could not estimate the warm \\h2 column densities. The {\\em Infrared Space Observatory} (ISO; Kessler et al. \\cite{kessler}), has allowed  us, for the first time, to measure directly the total column density of warm gas by observing pure-rotational  lines of \\h2. These trace gas with temperatures of a few hundreds Kelvin. ISO has also allowed us to study the \\h2 ortho-to-para ratio (OTPR), which can help determine       the possible heating mechanism and the origin of this molecule. Before  ISO, the \\h2 OTPR had been studied in regions with temperatures of $\\sim$ 2000~K, using the vibrational lines. In such shock-excited sources, one measures an OTPR of $\\sim$ 3 (Smith et al. \\cite{smith}), which is the local thermodynamical equilibrium (LTE) value for $T \\gsim$ 200 K. In contrast, for    regions heated mainly by ultraviolet (UV) radiation (Photodissociation regions [PDRs]), the vibrational lines give     OTPRs in the range of 1.2-2 (see e.g. Chrysostomou et al. \\cite{chrysostomou}). However, these low OTPRs  might  not be a consequence of an actual non-LTE ortho-to-para {\\em abundances} ratio but a result of optical depth effects in the fluorescence-pumping of the ortho-\\h2 (Sternberg \\& Neufeld \\cite{sternberg}). Using these considerations, one can  explain why the PDR in S140 exhibits an OTPR $\\sim$ 2 in the vibrational states but  3 in the lowest rotational levels.  There are two cases of non-equilibrium OTPR measured from the pure-rotational lines: the shock excited source HH54 (Neufeld et al. \\cite{neufeld}) and the PDR associated with the reflection nebula NGC 7023 (Fuente et al. \\cite{fuente99}). The first case has been explained using     the shocks model of Timmermann (\\cite{timmermann}), which involves transient heating by low velocity shocks. To explain  the non-LTE OTPR in NGC 7023, it was necessary to invoke a dynamic dissociation front.  To investigate the thermal balance of the GC clouds we have selected 18 clouds from the samples of H\\\"uttemeister et al. (\\cite{huttemeister93}) and Mart\\I n-Pintado et al. (\\cite{mp97}) and we have observed them with the ISO satellite. In this paper we present \\h2 observations toward two sources which show similar characteristics (also shown in the NH$_3$ studies of H\\\"uttemeister et al. \\cite{huttemeister93}), indicating that their heating mechanisms are also very similar. In particular, they show a non-LTE OTPR. The detection of OTPRs out of equilibrium in the GC clouds gives us new insights into the heating mechanism, since the gas must be heated to  several hundreds K almost without changing the OTPR of cold gas.   In Sects. 2 and 3,  we present observations and results, respectively, and in Sect. 4 we discuss the possible heating mechanism and the origin of the non-equilibrium OTPR.   \\begin{figure*}[p] \\centerline{\\psfig{figure=fig1.ps,bbllx=69pt,bblly=130pt,bburx=490pt,bbury=753pt,height=18cm}} \\caption[]{ Spectra of the two sources: {\\bf a-b} \\h2  spectra   taken with the SWS. {\\bf c-d} IRAM-30m spectra of the \\c18o(1--0), \\m13co(1--0), and \\c18o(2--1) lines. {\\bf e-f} LWS full grating spectra. Note the different radial velocities ranges in Fig. 1a-b and Fig. 1c-d. } \\label{figobs} \\end{figure*} ",
        "conclusions": "We have presented ISO SWS observations of the S(0), S(1), S(2), and S(3) pure-rotational lines of \\h2 and LWS observations of the dust continuum and IRAM-30m \\m13co and \\c18o observations  toward the GC molecular clouds \\muno and \\mtres. Using the CO data and dust column densities from the LWS spectra we estimate $\\sim$ 20 mag of visual extinction toward these sources. The two estimates   of the extinction agree within a factor of 2 for the standard gas-to-dust conversion factor. According to   the two components scenario proposed by H\\\"uttemeister et al. (\\cite{huttemeister93}), the low-J CO emission  arises from the cold gas component and is coupled to the dust  at a temperature of $<$ 30 K. The warm component  ($T \\sim$ 250 K) column density observed directly in \\h2 is, at least,  $\\sim$ 15 $\\%$ of the cold one and would have very little warm dust associated with it. From  the LWS spectra we set a conservative upper limit to the warm gas column density associated with the warm dust of 2~10$^{-3}$ times that of the warm \\h2 column density.  After correcting for the dust extinction, we  derive an OTPR of 1.0 $\\pm$ 0.4, which is far from the LTE value expected for the gas temperatures of 250 K. We have also compared the warm \\h2 column densities to the \\NH3 observations by H\\\"uttemeister et al. (\\cite{huttemeister93}), and derived \\NH3 abundances of $\\sim$ 2~10$^{-7}$, similar to those in the cold component (H\\\"uttemeister et al. \\cite{huttemeister93}).  The low dust temperatures, the high \\NH3 abundances, the large CO linewidths, the non-LTE \\h2 OTPR, in addition to the high gas temperatures suggest that the warm gas component is  heated  by  low velocity shocks with speeds of $\\sim$ 10 \\kms. To explain the OTPR we propose the following scenario. \\h2 is formed in the grain surfaces and ejected to gas phase with OTPR\\lsim 3. After $\\sim 10^6$ years, the time required to reach      the preshock OTPR=1, a low velocity shock heated the gas to  the observed temperatures of 250 K, but the OTPR was  almost unaltered because the timescale for the passage of such a shock is  much shorter than the ortho-to-para conversion timescale. Taking into account the shock timescale this occurred less than 10$^4$ yr ago. It is interesting to note that the timescale of the cloud's galactic rotation period is also $\\sim 10^6$ years. This fact suggests that the origin of the shocks can be related to large scale dynamics of the GC region."
    },
    "0002/astro-ph0002222_arXiv.txt": {
        "abstract": "Diluted arrays of many  optical apertures will be able to provide high-resolution snapshot images if the beams are combined according to the densified-pupil scheme. We show that the same principle can also provide coronagraphic images, for detecting faint sources near a bright unresolved one. Recent refinements of coronagraphic techniques, i.e. the use of a phase mask, active apodization and dark-speckle analysis, are also applicable for enhanced contrast. Implemented in the form of a proposed 50-500m Exo-Earth Discoverer array in space, the principle can serve to detect Earth-like exo-planets in the infra-red. It can also provide images of faint nebulosity near stars, active galactic nuclei and quasars. Calculations indicate that exo-planets are detectable amidst the zodiacal and exo-zodiacal emission faster than with a Bracewell array of equivalent area, a consequence of the spatial selectivity in the image. \\\\ {\\it Keywords}: extrasolar planets; instrumentation; image processing. ",
        "introduction": "\\label{intro} An early proposal for searching exo-planets involved coronagraphy aboard the \"Large Space Telescope\" (Bonneau {\\it et al.} 1975) which became the Hubble Space Telescope (HST). It was shown that telescope rotation and synchronous image subtraction could remove much of the coronagraphic residue, a speckled halo of star light expected to outshine the planet image by 3 or 4 orders of magnitude. The corresponding Faint Object Camera (FOC) was built by the European Space Agency, but the COSTAR corrector subsequently installed aboard HST to correct the figuring error of its main mirror destroyed the size matching of the Lyot stop in the FOC coronagraph and made it inoperable. Also, the aperture size had been shrunk from 3m to 2.4m, and the telescope rotation proved difficult to achieve. The FOC could however have succeeded since exo-planets brighter than expected were recently found by radial velocity measurements (Mayor and Queloz 1995, Marcy and Butler 1996, Delfosse {\\it et al.} 1998).\\\\  The rotation modulation principle was later adopted, together with a form of coronagraphic field-selective attenuation, by Bracewell and McPhie (1979) for their proposed interferometer exploiting the improved contrast of a planet in the  infrared. The concept has been refined by L\\'eger {\\it et al.} (1996) for the DARWIN proposal, and also adopted by Angel and Woolf (1997) for their Terrestrial Planet Finder. Both instruments are intended to observe in the 10$\\mu$m infra-red. At this wavelength, the Earth's luminosity is quoted to be $7.10^{6}$ times, or 17 magnitudes, fainter than the Sun's. \\\\  Coronagraphic imaging is now also considered for the Next Generation Space Telescope (Gezari {\\it et al.} 1997, Moutou {\\it et al.} 1998, Trauger {\\it et al.} 1998, Le F\\`evre {\\it et al.} 1998). The added refinement of \"dark-speckle imaging\" can further improve the detection sensitivity 10 to 1000 times (Boccaletti {\\it et al.} 1998a), while the use of a phase mask and the active attenuation of residual starlight in the annular field push the sensitivity towards the exo-planet detection threshold at visible wavelengths.\\\\  In this article we analyze and confirm the suggestion (Labeyrie 1999b) that these techniques of apodized imaging are also usable with large multi-aperture interferometric arrays operating in the densified-pupil imaging mode. Indeed, it was recently shown that such arrays  operated in this mode can produce snapshot images in a narrow field (Labeyrie 1996). We have studied the coronagraphic schemes applicable to such imaging conditions and find them suitable for exo-planet observing at infra-red wavelengths. At visible wavelengths they can provide higher angular resolution than  NGST coronagraphy,  allowing observations of fast orbiting exo-planets.\\\\  Our numerical simulations of imaging with the \"Exo-Earth Discoverer\" (EED) sketched in Fig. \\ref{eedschema} verify the expected imaging performance. The EED is a 36-element space interferometer with size $D$ of the order of 100m. Having free-flying telescopes (Labeyrie 1985, B\\'ely {\\it et al.} 1996), it can be considered as a possible precursor of much larger instruments  such as the  150 km Exo-Earth Imager proposed for the longer-term goal of making exo-planet portraits (Labeyrie 1999b). ",
        "conclusions": "The prospect for coronagraphic imaging with large interferometric arrays, at infrared and visible wavelengths, opens vast areas of science. The simulations presented show that searches for exoplanets at 10$\\mu$m  should be feasible once  the  formation flight of mirror elements in space becomes mastered. We recall the main conclusions : \\begin{itemize} \\item{ Pupil densification allows direct imaging and coronagraphic nulling.} \\item{ A phase-mask, replacing the classical Lyot opaque mask, cleans the image closer to the attenuated star, and is more tolerant to the unavoidable narrow gaps between sub-pupils in the densified pupil.  Both achromatisation of the mask and low-resolution spectroscopy ($\\lambda / \\Delta \\lambda = 20$) are desirable for analysing  the images obtained.} \\item{At given collecting area, more sub-apertures is better in terms of image field and signal/noise ratio. } \\item{The optimal wavelength is a trade-off between  resolution, stellar spill-off, planet contrast, background emission and detector performance. At visible wavelengths, a monolithic 8m telescope often resolves the planet from its star and then has the same planet-searching efficiency as an EED of identical area.} \\item{ In addition to planet searches, direct imaging with the EED is  also of interest for the other science goals briefly described above.} \\item{With respect to Bracewell's detection scheme, and its DARWIN and TPF proposed implementations,  the image formation in the EED improves markedly, at equal collecting area, the discrimination of  an exo-planet amidst the  zodiacal and exo-zodiacal nebulosity.  The EED however requires 36 apertures  at least, but these are smaller and can be implemented with identical mirror segments compatible with the nano-satellite philosophy }.  \\end{itemize} The coronagraphic image is compatible with field spectrography.  Additional simulations are needed to compare the possible exposure routines, assess the desirability of a fine wavefront sensor, and refine the noise models. More work is also needed to design optimal optical trains, including the coronagraphic device and its spectro-imaging attachment. \\\\  Acknowledgments : Jean-Marie Mariotti his constructive was among the first to understand our ideas and on densified-pupil imaging encouraged us to explore their application. His premature death deprived us of his insight and critical appreciation, particularly regarding these results.  \\newpage   \\Large{\\bf{Figure Captions}}\\\\[1cm] \\normalsize  Figure \\ref{eedschema}: Possible optical train for the Exo-Earth Discoverer. Mirror elements M1 focus starlight in the focal combiner C (detail below), on field mirrrors FM, which form a densified image of the pupil on mirror segments PM. A combined image is focused in F at the entrance of a coronagraph Co. Mirrors PM are carried by actuators for tip-tilt and piston corrections.\\\\  Figure \\ref{pupil}: Interferometer apertures (top row) with disk (left) or hexagon (right) subpupils, and densified version of them obtained in the exit pupil (second row). The 100 to 200m Exo-Earth-Discoverer has a sparse entrance aperture, but a fully densified exit pupil such as shown at right in the second row. The star image (third row), with its diffractive envelope, has most energy in the central Airy peak, as seen in the intensity profiles (bottom).\\\\  Figure \\ref{maskh}: Top: Central darkening versus phase-mask size with a fully densified pupil of 36 adjacent hexagons, measured at the maximum ( dotted line) and across the area  (solid line) of the Airy peak . The  optimal mask size found is $(0.555 \\pm 0.005)\\lambda/D$. Bottom: Azimuthally averaged and normalized profiles of the unmasked Airy pattern (solid line), coronagraphic pattern (dashed line) and residual scattered pattern after subtraction of  the reference frame (dotted line), including photon and readout noise.  Both long exposures incorporated 20 different maps of ($\\lambda/170$ rms) wavefront bumpiness. The radii of the first dark ring in the unmasked ($\\rho_u$), the masked ($\\rho_m$) and the sub-aperture ($\\rho_f$) images are indicated. Also indicated is the level of Venus (dash-dot), Earth (dash-dot-dot-dot) and Mars (long dashes) multiplied by the sub-aperture's diffraction function .\\\\  Figure \\ref{ecart}: Coronagraphic attenuation of star light versus exit-pupil density. The pupil density is measured as the relative area of the sub- pupils in the exit pupil, while the attenuation is measured in the annular field ($12\\lambda/D$) and normalized to the intensity of the unmasked peak. The dotted curve indicates that the mask size was re-optimized for each density value, while  the solid curve  indicates that the initial value ($0.555\\lambda/D$) was maintained . Numbers indicate the sizes of the re-optimized masks in units of $\\lambda/D$.\\\\  Figure \\ref{result}: Coronagraphic image of a  solar system twin observed at 20pc ($m_v=6,33$) with the Sun attenuated to detect Venus, Earth and Mars. The magnitude differences at 10 $\\mu$m are respectively 16.4, 17.1, 19.8 relative to the parent star. The simulated array has 36 telescopes of 0.6m, providing a collecting area similar to a DARWIN (40m$^2$). To discriminate the dispersed secondary peaks from the white primary peaks, the final image was obtained with a combination of 3 wavelengths (9.5, 10 and 10.5$\\mu m$).\\\\ In (a), the star is removed to show the Venus, Earth and Mars components of the image, with their primary (V, E, M) and secondary ($V_1$, $V_2$, $E_1$, $E_2$, $M_1$, $M_2$) peaks. The peripheral attenuation by the diffraction envelope makes the secondary peaks of Mars brighter than the primary peak, since they are closer to the field center. No photon noise is present here.\\\\ The separations are $1.75\\lambda/D$, $2.43\\lambda/D$ and $3.69\\lambda/D$, corresponding, for a 100m array, to 36, 50 and 76 milli-arcsecond at 10 $\\mu$m, respectively. A circle indicates the  field limit, at the first dark ring of the diffraction envelope of a single 0.6m-aperture. On the sky, the corresponding  field diameter is $12\\lambda/D$, amounting to 0.24 arc-second at 10 $\\mu$m with a 100m array of 36 elements.\\\\ Image (b) shows a 10 hours exposure with subtraction of a reference star image similarly obtained. Photon noise and readout noise (100 e-/pix/frame) were included. A perfect wavefront was assumed. Image (c) is similar, but with 59 nm RMS wavefront bumpiness ($\\lambda /170$), of 10mn lifetime. Two different sets of 60 phase maps served for the object and the reference star. This level of bumpiness has little effect in the usable field, within the white circle. A further gain in sensitivity is achievable in the dark-speckle mode.\\\\   \\newpage"
    },
    "0002/astro-ph0002258_arXiv.txt": {
        "abstract": "With high sensitivity imaging arrays, accurate calibration is essential to achieve the limits of detection of space observatories. One can simultaneously extract information about the scene being observed and the calibration properties of the detector and imaging system from redundant dithered images of a scene. There are large differences in the effectiveness of dithering strategies for allowing the separation of detector properties from sky brightness measurements. In this paper, we quantify these differences by developing a figure of merit (FOM) for dithering procedures based on their usefulness for allowing calibration on all spatial scales. The figure of merit measures how well the gain characteristics of the detector are encoded in the measurements, and is independent of the techniques used to analyze the data. Patterns similar to the antenna arrangements of radio interferometers with good $u-v$ plane coverage, are found to have good figures of merit. We present patterns for both deep surveys of limited sky areas and for shallow surveys. By choosing a strategy that encodes the calibration in the observations in an easily extractable way, we enhance our ability to calibrate our detector systems and to reach the ultimate limits of sensitivity which are required to achieve the promise of many missions. ",
        "introduction": "In order to achieve their required performance, many observing systems must observe with sensitivities near their confusion limits. Many instruments are capable of reaching these limits in crowded stellar fields such as the Galactic center. Future instruments such as the Infrared Array Camera (IRAC) and the Multiband Imaging Photometer (MIPS) on the {\\it Space Infrared Telescope Facility} ({\\it SIRTF}) and those planned for the Next Generation Space Telescope (NGST) will be able to reach limits in which the confusion of extragalactic sources becomes significant. In general, the measurement noise is determined by both the statistical fluctuations of the photon flux and uncertainties in detector gain and offset. Any successful calibration procedure must determine these detector parameters sufficiently accurately so that their uncertainties make small contributions to the measurements errors compared to those of the background fluctuations. If the science done with the instrument requires substantial spatial or temporal modeling, calibration requirements become more demanding, ultimately requiring similar integration time for observation and calibration as in the case of the {\\it COBE} FIRAS instrument (Mather et al. 1994; Fixsen et al. 1994). Additionally, in such cases robust error estimators are often needed. A common method to determine the instrument calibration is to look at known calibration scenes (e.g. a dark shutter, an illuminated screen, or a blank region of sky) of different brightnesses to deduce gain and offset of each detector pixel. This requires a well characterized calibration source and often a change in instrument mode to carry out the measurement. This procedure may introduce systematic errors relating to the extrapolations from the time and conditions of the calibration observations to the time and conditions of the sky observations and from the intensity (and assumed flatness) of the calibration source to the intensity of the observed sky. A different approach is to use the measurements of the sky alone to extract the calibration data for the system. By using the sky observations for calibration, the systematic errors introduced by applying a calibration derived from a distinctly different data set are eliminated. Such methods require a set of dithered images, where a single sky location is imaged on many different detector pixels.  Typical CCD and IR array data reduction procedures for a set of dithered images make use of a known or measured dark frame ($F^p$) and derive the flat field ($G^p$) through taking the weighted average or median value of all data ($D^i$) observed by each detector pixel $p$ ($i \\in p$) in a stack of dithered images (e.g. Tyson 1986, Tyson \\& Seitzer 1988, Joyce 1992, Gardner 1995). The least squares solution of \\begin{equation} {\\cal D}^i = G^p S^0 \\end{equation} where ${\\cal D}^i = D^i - F^p$ ($i \\in p$) and $S^0$ is the perfectly flat sky intensity, for $G^p$, the flat field, is \\begin{equation} G^p = \\frac{\\sum_{i \\in p}{\\cal D}^i W_i}{\\sum_{i \\in p}W_i} \\frac{1}{S^0} \\end{equation} which is simply the weighted average of the data collected by each detector pixel normalized by the constant sky intensity (to be determined later though the absolute calibration of the data). The weights, $W_i$, are normally determined by the inverse variance of the data, but may also be set to zero to exclude sources above the background level. The use of the median, instead of the weighted average, also rejects the outliers arising from the observations of real sources instead of the flat background, $S^0$, and formally corresponds to a minimization of the mean absolute deviation rather than a least squares procedure. In either form, this method requires observations of relatively empty fields where variations in the background sky level are not larger than the faintest signal that is sought. Thus, throughout this paper we refer to such procedures as ``flat sky'' techniques. As instrumentation improves and telescope sensitivity increases, this condition is becoming harder to fulfill. In fields at low Galactic latitude, stellar and nebular confusion can be unavoidable, and at high latitude deep imaging (particularly in the infrared) is expected to reach the extragalactic confusion limit. In such cases, because of the complex background, and in other cases where external influences (e.g. moonlight, zodiacal light) create a sky background with a gradient, the flat sky approach does not work and a more comprehensive approach must be used.  Such an approach has been presented by Fixsen et al. (2000) who describe the general least squares solution for deriving the sky intensity $S^{\\alpha}$ at each pixel $\\alpha$, in addition to the detector gain (or flat field) $G^p$ and offset (or dark current + bias) $F^p$ at each detector pixel $p$, where each measurement, $D^i$, is represented by \\begin{equation} D^i = G^p S^{\\alpha} + F^p. \\end{equation} (Throughout this paper we refer to the procedure described by Fixsen et al. (2000) as the ``least squares'' procedure.) They show how the problem of inverting large matrices can be circumvented, and how the formulation of the problem allows for explicit tracking of the uncertainties and correlations in the derived $G^p$, $F^p$, and $S^{\\alpha}$. Fixsen et al. also show that although the formal size of the matrices used in the least squares solution increases as $P^2$, where $P$ is the number of pixels in the detector array, the number of non-zero elements in these matrices increases only as $M\\times P$, where $M$ is the number of images in the data set. In practice, the portion of the least squares solution for the detector gains and offsets is calculated first, and then the data are corrected to produce images of the sky ($S^{\\alpha}$) that are registered and mapped into a final single image. Because this approach explicitly assumes a different sky intensity at each pixel, the crowded or confused fields that can cause the flat sky technique to fail are an aid to finding the least squares solution. Thus, the need for chopping away from a complex source in order to observe a blank sky region is eliminated. The simultaneous solution for both the detector gain and offset also eliminates the need for dark frame measurements, although if dark frame measurements are available then they can be used with the other data to reduce the uncertainty of the procedure. We note that this general least squares approach may also be applied in non-astronomical situations (e.g. terrestrial observing) where complex images are the norm.  The flat sky technique works well in situations where all detector pixels spend most of the time observing the same celestial calibration source, namely the flat sky background. For this technique, dithering is required only to ensure that all pixels usually do see the background. Because all pixels have observed the same source, the relative calibrations of any two pixels in the detector are tightly constrained, regardless of the separation between the pixels, i.e. \\begin{equation} \\frac{G^1}{G^2} = \\frac{G^1 S^0}{G^2 S^0} = \\frac{{\\cal D}^1}{{\\cal D}^2}. \\end{equation} However, in the more general least squares solution of Fixsen et al. (2000), each sky pixel ($S^{\\alpha}$) represents a different celestial calibration source. The only pixels for which the relative calibrations are tightly constrained are those that through dithering have observed common sky pixels. Pixels that do not observe a common sky pixel are still constrained, though less directly, by intermediate detector pixels that do observe common sky pixels. For example, the relative calibration of detector pixels 1 and 3 which observe sky pixels $\\alpha$ and $\\beta$ respectively, but no common sky pixels, may be established if an intermediate detector pixel 2 does observe both sky pixels $\\alpha$ and $\\beta$, i.e. \\begin{equation} \\frac{G^1}{G^3} = \\frac{G^1 S^{\\alpha}}{G^2 S^{\\alpha}} \\frac{G^2 S^{\\beta}}{G^3 S^{\\beta}} = \\frac{{\\cal D}^{1\\alpha}}{{\\cal D}^{2\\alpha}} \\frac{{\\cal D}^{2\\beta}}{{\\cal D}^{3\\beta}}. \\end{equation} Other detector pixels might require multiple intermediate pixels to establish a relative calibrations. As the chain of intermediate pixels grows longer, the uncertainty of the relative calibration of the two pixels also grows. Therefore, when applying the least squares solution, the exact dither pattern becomes much more important than in the flat sky technique. For the least squares solution to produce the smallest uncertainty, the dither pattern should be one that establishes the tightest correlations between all pairs of detector pixels using a small number of dithered images. Even if one is only interested in small scale structure on the sky (e.g. point sources), it is still important to have the detector properly calibrated on {\\it all} spatial scales to prevent large scale detector variations from biasing results derived for both sources and backgrounds imaged in different parts of the array.  Whether obtained by flat sky, least squares, or other techniques, the quality of the calibration is ultimately determined by its uncertainties. For the least squares solution of Fixsen et al. (2000), understanding the uncertainties is relatively straight forward, because it is a linear process, i.e. $P^{\\alpha} = L^{\\alpha}_i D^i$ where $P^{\\alpha}$ is the set of fitted parameters, $D^i$ is the data, and $L^{\\alpha}_i$ is a linear operator. Then, given a covariance matrix of the data, $\\Sigma^{ij}$, the solution covariance matrix is $V^{\\alpha\\beta} = L^{\\alpha}_i L^{\\beta}_j \\Sigma^{ij}$. For a nonlinear process such as a median filter the uncertainties are harder to calculate. The diagonal terms of the covariance matrix of the solution might be sufficiently well approximated by Monte Carlo methods, but the off-diagonal components are far more numerous and often more pernicious as the effects can be more subtle than the simple uncertainty implied by the diagonal components. For this reason the off-diagonal components are often ignored. Creating final images at subpixel resolution (e.g. ``drizzle'', Fruchter \\& Hook 1998) may introduce additional correlations beyond those described by the covariance matrix, and disproportionately increase the effects of the off-diagonal elements of the correlation matrix. Accurate knowledge of all these uncertainties is especially important for studies that seek spatial correlations within large samples, such as deep galaxy surveys or studies of cosmic backgrounds, so that any detected correlations are certifiably real and not artifacts caused by the calibration errors and unrecognized because of incomplete or faulty knowledge of the uncertainties.  Table 1 itemizes some of the features of each data analysis technique. The remainder of this paper is concerned with characterizing what makes a dither pattern good for self-calibration purposes using the least squares solution. We present a ``figure of merit'' (FOM) which can be used as a quantitative means of ranking the suitability of different dither patterns (\\S2). We then present several examples of good, fair and poor dither patterns (\\S3), and investigate how changes to the patterns affect their FOM. In \\S4, we show how dithered data can be collected in the context of both deep and shallow surveys. We also investigate the combined effects of dithering and the survey grid geometry on the completeness of coverage provided by the survey. Section 5 discusses miscellaneous details of the application and implementation of dithering. Section 6 summarizes the results. ",
        "conclusions": "We have shown that proper selection of observing strategies can dramatically improve the quality of self-calibration of imaging detectors. We have established a figure of merit (Eq. 6) for quantitatively ranking different dither patterns, and have identified several patterns that enable good self-calibration of a detector on all spatial scales. The layouts of radio interferometers correspond to good dither patterns. Both the highly ordered Reuleaux triangle pattern and the unstructured random pattern provide good FOM with moderate or deep observations. This indicates that good patterns must sample a range of spatial scales without redundancy, and if this condition is met, then secondary characteristics of the patterns or instrument constraints may determine the actual choice of the dither pattern. Any dither pattern must contain steps as large as half the size of the detector array if large scale correlations are to be effectively encoded in the dithered data set. Deep surveys can take advantage of the use of a single good dither pattern. Shallow surveys can obtain good FOM by altering the dithers used at each of the survey grid points. Using a fixed pattern throughout a shallow survey makes it difficult or impossible to apply a self-calibration procedure to the resulting data sets. The use of triangular instead of square survey grids can be more efficient in executing complete-coverage surveys when the array orientation cannot be set to match the survey grid. Good dither patterns and survey strategies can be devised even in some seemingly restricted situations. The ultimate importance of dithering and a good FOM will depend on the nature of the instrument and the data and on the scientific goals. For many goals, obtaining a larger quantity of data may not be an adequate substitute for obtaining data with a good FOM."
    },
    "0002/astro-ph0002402_arXiv.txt": {
        "abstract": "Photons with very high energy up to TeV (VHE) emitted from active galactic nuclei (AGNs) provide some invaluable information of the origin of $\\gam$-ray emission. Although 66 blazars have been detected by {\\it EGRET}, only three low redshift X-ray selected BL Lacs (Mrk 421, Mrk 501, and 1ES 2344+514) are conclusive TeV emitters (PKS 2155-304 is a potential TeV emitter) since VHE photons may be absorbed by cosmological background infrared photons ({\\it external} absorption). Based on the ``mirror'' effect of clouds in broad line region, we argue that there is an {\\it intrinsic} mechanism for the deficiency of TeV emission in blazars. Employing the observable quantities we derive the pair production optical depth $\\tgg(\\eps_{\\rm obs})$ due to the interaction of VHE photons with the reflected synchrotron photons by electron Thomson scattering in broad line region. This sets a more strong constraints on very high energy emission, and provides a sensitive upper limit of Doppler factor of the relativistic bulk motion. It has been suggested to distinguish the {\\it intrinsic} absorption from the {\\it external} by the observation on variation of multiwavelegenth continuum. ",
        "introduction": "Clearly, the high energy $\\gam$-ray emission is an important piece in the blazar puzzle because the $\\gam$-ray observations of blazars provide a new probe of dense radiation field released through accretion onto a supermassive black hole in the central engine (Bregman 1990). The Energetic Gamma Ray Experiment Telescope ({\\it EGRET}) which works in the 0.1--10GeV energy domain has now detected and identified 66 extragalactic sources in 3th catalog (Mukherjee et al 1999). All these objects are blazar-type AGNs whose relativistic jets are assumed to be close to the line of sight to the observer. It seems unambiguous that the intense gamma-ray emission is related with highly relativistic jet.  It has been generally accepted that the luminous gamma-ray emission is radiated from inverse Compton, but the problem of seed photons remains open for debate. The following arguments have been proposed: (1) synchrotron photons in jet (inhomogeneous model of synchrotron self Compton) (Maraschi, Ghisellini \\& Celotti 1992); (2) optical and ultraviolet photons directly from the accretion disk (Dermer \\& Schlikeiser 1993); (3) diffusive photons in broad line region (BLR) (Sikora, Begelman \\& Rees 1994, Blandford \\& Levinson 1995); (4) the reflected synchrotron photons by electron mirror in broad line region, namely, the reflected synchrotron inverse Compton (RSC) (Ghisellini \\& Madau 1996). These mechanisms may operate in different kinds of objects, however there is not yet a consensus on how these mechanisms work. Also it is not clear where the $\\gam$-ray emission is taking place largely because of uncertainties of soft radiation field in the central engine.  On the other hand, VHE observations (Kerrick et al 1995, Chadwick et al 1999, Roberts et al 1999, Aharonian et al 1999) are making attempts to explore the radiation mechanism because they may provide some restrictive constraints (Begelman, Rees \\& Sikora 1994, Mastichiadis \\& Kirk 1997, Tavecchio, Maraschi \\& Ghisellini 1998, Coppi \\& Aharonian 1999, Harwit, Protheroe and Biermann 1999). Based on the simple version of SSC model, Stecker, de Jager \\& Salamon (1996) predicted a large number of low redshift X-ray selected BL Lacs as TeV candidates, taking into account that the presence of intergalactic infrared radiation field including cosmic background leads to strong absorption of TeV photons from cosmological emitters (Stecker \\& de Jager 1998). It is suggested to form an extended pair halo in cosmological distance due to the {\\it external} absorption (Aharonian, Coppi, \\& Voelk 1994). However, so far only three X-ray selected BL Lacs have been found to be TeV emitters by Whipple telescope ($E>300$GeV), in addition, photons higher than 0.3TeV in the X-ray-selected PKS 2155-304 with redshift $z=0.116$ has been detected photons 0.3TeV by Durham Mrk 6 telescope (Chadwick et al 1999). The recent measurements of intergalactic infrared field is quite different from the previous observations (Madau et al 1998, Steidel 1998).  Although this {\\it external} absorption is definitely important, the critical redshift $z_{c}$ beyond which cosmological back ground radiation and intergalactic infrared fields will absorb VHE photons remains uncertain. Especially the recent VHE observations show that Mrk 501 emits 25 TeV photons (Aharonian et al 1999). Evidently this suggests that the {\\it external} absorption can not efficiently attenuate the VHE photons from reaching us across distances of 100 Mpc. It is highly desired to accurately probe the star formation rate in order to determine the critical redshift $z_{c}$.  Thus it seems significant to study the {\\it intrinsic} mechanism for the deficiency of TeV photons from $\\gamma$-ray loud AGNs disregarding the absorption by intergalactic infrared radiation field. A larger Lorentz factor of the jet implies higher density of the external photons in the blob, if the reflection of clouds in broad line region works, and therefore stronger absorption of high energy $\\gamma$-rays (Celotti, Fabian \\& Rees 1998). Here we argue based on the hypothesis of Ghisellini \\& Madau (1996) that the energy density of reflected synchrotron photon is high enough for pair production via interaction of gamma-ray photons by inverse Compton scattering with reflected synchrotron photons if the bulk velocity is high enough. Further we apply the present constraint to the representative individual objects, Mrk 421 and 3C 279. ",
        "conclusions": "  1) The parameters $\\lrsc$ and Doppler factor $\\cd$ in FSRQs are systematically greater than that in BL Lacs. This will cause the more stronger ``{\\it intrinsic}'' absorption of VHE photons in FSRQs than that in BL Lacs. It is predicted that there is general absence of very high energy emission in FSRQs, owning to the attenuation of VHE photons by the BLR reflection of synchrotron emission.  2) The mirror model provides a new constraint on relativistic bulk motion. That {\\it intrinsic} absorption of TeV photons may operate in some objects, especially in FSRQs. This constraint is cause by the motion of blob itself.  Although the origin of $\\gam$-ray emission in blazars still remains open, VHE observations sets strong constraints on blazar's radiation mechanism. These constraints are: (1) brightness temperature exceeding the Kellermann-Pauliny-Toth (Begelman et al 1994), (2) multiwavelength light curves based on the homogeneous model (Mastichiadis \\& Kirk 1997), (3) high energy variations in X-ray and $\\gam$-ray including interaction with background IR radiation (Coppi \\& Aharonian 1999). These constraints are mainly based on SSC model. The deficiency of VHE photons from high redshift $\\gamma$-ray loud blazars may be explained by the interaction of the cosmic background radiation fields with the VHE photons. However the possible alternative mechanism may be due to the {\\it intrinsic} attenuation by the reflected synchrotron photons. Three BL Lac objects have been found to show $H\\alpha$ and $H\\beta$ emission, indicating the existence of broad line region in these objects (Vermeulen et al 1995, Corbett et al 1996), even Mrk 421 has been detected weak luminosity of a broad emission line (Morgani, Ulrich \\& Tadhunter 1992). The increasing evidence of the presence of broad emission lines in BL Lacs lends the possibility that the reflected external photons might be the main source of seed photon in this kind of blazars. Distinguishing the two different mechanisms might be traced by the following-up observations in other wavebands because a pair cascade process may be developed, forming a pair halos in the {\\it external} absorption (Aharonian, Coppi \\& Voelk 1994). Such an extended halo due to {\\it external} absorption may be of very long variable timescale at least $\\sim 10^3$ yr (corresponding to one mean free path) [see equation (24)]. However if  the {\\it intrinsic} absorption works in the central engine, the time-dependent synchrotron self-Compton model including pair cascade (Wang, Zhou \\& Cheng 2000) could predict the interesting spectrum and light curves, which may interpret the variations of PKS 2155-304 (Urry et al 1997). The other radiative properties of such an extended pair halo are needed to be studied in order to distinguish the {\\it intrinsic} absorption from the {\\it external} one."
    },
    "0002/astro-ph0002416_arXiv.txt": {
        "abstract": "We develop an algorithm for the reconstruction of the two-dimensional mass distribution of a gravitational lens from the observable distortion of background galaxies. From the measured reduced shear $\\gamma_i/(1-\\kappa)$ the lens mapping is obtained, from which a mass distribution is derived. This is unlike other methods where the convergence $\\kappa$ is directly obtained.  We show that this method works best for sub-critical lenses, but can be applied to a critical lens away from the critical lines.  For finite fields the usual mass-sheet degeneracy is shown to exist in this method as well.  We show that the algorithm reproduces the mass distribution within acceptable limits when applied to simulated noisy data. ",
        "introduction": "\\label{sec:intro}  The mapping of the distribution of matter in extended regions around rich clusters of galaxies, from the systematic distortions of background galaxies due to gravitational lensing, has increasingly become feasible and popular over the last decade.  This \\emph{shear} is measured from the quadrupole moments of the images of the background galaxies in terms of the local `polarization' of an image compared to its assumed intrinsic circular form.  In order to remove the effect of the source ellipticity distribution, the parameters are averaged over a large number of source galaxies.  The two-dimensional mass distribution of the lens is reconstructed from the shear map away from the critical lines (`weak reconstruction'). Following the pioneering work of Tyson \\etal\\ (1990), the idea was quantitatively developed by Kaiser and Squires (1993, also Kaiser 1995 and Squires and Kaiser 1996) and Seitz and Schneider (1995, 1996). Since then, several interesting variants on this theme have appeared in the literature, \\eg\\ using maximum likelihood cluster reconstruction (Bartelmann \\etal\\ 1996, Seitz \\etal\\ 1998), or using methods based on the variational principle (Lombardi and Bertin 1999) or maximum entropy (Bridle \\etal\\ 1998).  Since the shear data contain complete information about the mass distribution in two independent fields $\\gamma_1 /(1-\\kappa)$ and $\\gamma_2/(1-\\kappa)$, it is possible to obtain several algorithms to estimate the mass distribution from the measured shear.  The observed shear is available only on a finite grid. Due to the unknown intrinsic distribution of the ellipticities of the source galaxies, and the effect of the distortion of the PSF due to observing conditions (seeing, tracking etc.), the measured shear is noisy.  It is therefore desirable to develop new algorithms in the hope that they might be able to deal with the noise better than other methods.  In this paper we develop an algorithm (LM: Lens Mapping algorithm) for the lens mass reconstruction from the measured reduced shear, $g_i = \\gamma_i/(1-\\kappa)$. The method involves two steps. First, the  lens mapping is reconstructed from the reduced shear, which can be done uniquely with the assumption that the lens mapping goes to identity far away from the lens.  As the second step, we show in \\S\\ref{sec:algorithm} that for a sub-critical lens, the surface mass density can be reconstructed completely from the derived lens mapping. In reality, however, the measured shear is available only in a finite region of the lens plane. In this case the lens mapping cannot be uniquely obtained and hence the LM algorithm exhibits a mass-sheet degeneracy, which is also present in other methods of mass reconstruction.  In \\S\\ref{sec:finite}, we characterize this degeneracy in the context of the LM algorithm. The performance of this method with discretely sampled, noisy data is dealt with in \\S\\ref{sec:noisy}, where we demonstrate the various features of the LM algorithm, by reconstructing the mass distribution for an analytically given shear field, and separately showing the effects of discrete sampling and of uniform noise (due to measurement error) on the input data.    \\begin{figure*} \\newlength{\\figwidth} \\setlength{\\figwidth}{\\textwidth} \\addtolength{\\figwidth}{-\\columnsep} \\setlength{\\figwidth}{0.5\\figwidth}  \\begin{minipage}[t]{\\figwidth} \\mbox{} \\vskip -7pt \\centerline{\\includegraphics[width=0.82\\linewidth,angle=0]{vecfield.eps}} \\end{minipage} \\hfill \\begin{minipage}[t]{\\figwidth} \\mbox{} \\vskip -7pt \\centerline{\\includegraphics[width=0.82\\linewidth,angle=0]{grid.ps}} \\end{minipage} \\caption{ [a] The coordinate system used in this paper.  We show the lens-plane images of two straight lines passing through ($y_1$, $y_2$) in the source plane; [b] An equally-spaced Cartesian grid (in the source plane) is lensed by a two-component mass model, both PICMDs with zero ellipticity and $k_0 = 0.8$ and $0.5$ respectively (described in \\S\\ref{sec:algorithm}). The axes are marked in units of the core radius $r_c$, assumed to be same for both the models. The separation between the centers of the two distributions is $10\\,r_c$.  \\label{fig:grid} }  \\end{figure*} ",
        "conclusions": "We have described a simple algorithm for mass reconstruction of a cluster lens by directly obtaining the lens mapping from the measured reduced shear and taking its derivatives to compute the surface mass density $\\kappa$.  The method works best for sub-critical lenses, where it can give the mass distribution at all points, but it can work for critical lenses as well in limited regions of the lens plane, away from the critical lines.  The algorithm is shown to have a mass-sheet degeneracy of the same type as exists in the other methods of reconstruction if the field in which the reduced shear is available is finite.  We have tested the algorithm on discretely sampled noisy (simulated) data and have found that it reproduces the mass distribution within acceptable limits."
    },
    "0002/astro-ph0002146_arXiv.txt": {
        "abstract": "We present near-IR photometry of a selected sample of southern hemisphere E+A galaxies. The sample includes 50 galaxies from nearby ($z \\sim 0.05$) and distant ($z \\sim 0.3$) clusters as well as E+A galaxies from the field ($z \\sim 0.1$). We also observed 13 normal early-type galaxies from the field and from clusters to be compared with the E+A sample. The photometry includes $J$, $H$ and $K_s$ apparent magnitudes and colors. Observed colors are obtained from the apparent total magnitudes and compared to spectrophotometric models of galaxy evolution GISSEL96. There is an overall agreement between integrated colors of models and observed ones, for both the E+A located in clusters and in the field, at $z \\lesssim 0.1$. However, large differences are found between colors predicted from models and those observed in E+A galaxies located in clusters at $z \\sim 0.3$.  We also compute rest-frame colors for all the galaxies using two different sets of K-corrections, and obtain average colors for all the samples.  We investigate systematic properties of the E+A sample as a function of their environment. Results seem to indicate that cluster E+As (at low redshift) are bluer than field E+As at $z \\sim 0.1$. Even this conclusion does not depend whether we use comoving or rest-frame colors, nor on the models used to obtain rest-frame colors; the difference is not significant enough, considering color dispersions between the samples. If differences are real, they could imply different stellar content for the E+A galaxies located in the field, compared to the cluster E+A. ",
        "introduction": "The attention drawn to E+A galaxies (or post-starbursts) has increased since it was claimed that their fraction observed in clusters is correlated with the Butcher-Oemler effect \\citep{butcher78, dressler83, rakos95}. The E+A galaxies present a peculiar spectrum in the optical: strong Balmer absorption lines, representative of a large population of A and B stars, but a {\\em lack} of emission lines typical of blue, star forming galaxies, like [OII]$\\lambda3727$, [OIII]$\\lambda$5007, and H$\\alpha$. Because of the further detection of metallic absorption lines such as Mg b $\\lambda 5175$, Ca H \\& K $\\lambda 3934$, 3968 and Fe $\\lambda5270$, indicative of an old population dominated by G, K and M spectral types they were called E+A \\citep{dressler83}. In addition, their spectra in the {\\em optical} cannot be reproduced simply by a young stellar component without nebular emission lines: it is necessary to add an old population of stars, like the one present in quiescent elliptical galaxies \\citep{liu96}.  During the last 3 years, research on the E+A galaxies has seen a revival after the discovery of more such galaxies not only in nearby clusters, like Coma and others (e.g. see \\citet{caldwell97} and references therein), but also in the field; in particular those discovered during the Las Campanas Redshift Survey (LCRS, Shectman \\etal 1996) by \\citet{zabludoff96}. An interesting feature is the apparent existence of two classes of E+A galaxies \\citep{couch87}. One is formed by ``blue'' post-starburst galaxies and the other by redder, H$\\delta$-strong (HDS) galaxies \\citep{fabricant91}. These subclasses of E+As have colors and absorption line features related to the morphology: HDS E+As have in general a noticeable bulge and/or spheroidal component when compared with the blue class. In addition, the HDS class can be divided into 2 subclasses: bulge and disk HDSs (as observed in A665 and in Coma \\citep{franx93}).  An almost unexplored domain of these E+A galaxies is their near-IR properties. In fact, no catalogue or systematic observations exist on this subject. Is the bright, red population, dominated mainly by giants and stars of the asymptotic giant branch (AGB) of the E+As, different from the red population of other elliptical galaxies, in particular the perturbed ellipticals? Is there any conspicuous signature in the near-IR colors of the E+A galaxies, as in the optical wavelengths? Do the cluster E+A galaxies have bluer colors than those in the field at these wavelengths? Are the near-IR colors of the E+A galaxies similar to the colors of normal galaxies, as predicted by spectrophotometric models?  In this paper, we investigate these questions with new data taken at Las Campanas Observatory, in Chile. The sample includes E+A galaxies from the field (most of them from the LCRS) and from nearby clusters as well as clusters at $z \\sim 0.3$.  The paper is organized as follows. In \\S 2 we present  the galaxies selected for this work. In \\S 3 we explain the observations, and in \\S 4 the data reduction procedures. In \\S 5 we present the results of the photometry, the apparent magnitudes, colors, and K-corrections. Both, {\\em observed}  and {\\em rest-frame} colors are compared with colors obtained from spectrophotometric models of galaxy evolution in \\S 6. In \\S 7 we discuss the limitations of our results and the implications for some properties of the E+A galaxies from our near-IR colors. Conclusions are presented in \\S 8. ",
        "conclusions": "The E+A galaxies reported here include 32 galaxies from clusters and 18 galaxies from the field. In addition, 13 nearby galaxies which do not present post-starburst activity, were observed (5 located in clusters at $z \\sim 0.05$ and 8 located in the field at very low redshift). All the galaxies have been observed in the near-IR bands $J$, $H$ and $K_s$ during photometric nights at Las Campanas Observatory. Total apparent magnitudes and colors were derived. The color-color diagram $H - K_s$/$J - H$  of the observed galaxies is compared to the expected corresponding colors of spectrophotometric models of galaxy evolution, at different redshifts. The models are those generated by GISSEL96 \\citep{charlot96}. There is an overall agreement between these expected colors and the observed ones, for the E+A located in nearby clusters ($<z> \\sim 0.05$) and for E+As located in the field ($<z> \\sim 0.1$). The comparison of the colors of these two samples shows that even though cluster E+As appear bluer than field E+As, the color difference is only significant at $\\sim 1.5\\sigma$ level, and therefore we cannot strongly affirm that stellar population differences are observed between these two populations.  The colors of the E+A galaxies located in more distant clusters $z = 0.3$, on the other hand, do not agree with the color expected from models. In the mean, they appear bluer than expected in $H - K_s$ (by $\\sim 0.3$mag) and redder in $J - H$ (by $\\sim 0.15$ mag). The possible interpretation of the failure is strong internal reddening (mostly in the $J$ band), not considered in models.  In order to derive a more complete comparison with models, rest-frame colors were also obtained using two different sources of K-corrections: one based on the work of \\citet{poggianti97}, and the other computed using the spectrophotometric model PEGASE \\citep{fioc97}. We have shown that such K-corrections can be significant for $z \\sim 0.2$ in the $K$ bands (or any band centered at 2$\\mu m$), although they are not a strong function of spectral type. In addition, large differences exist in the K-corrections between these two models, having a large impact on the derived quantities, like rest-frame colors for high redshift galaxies. We have compared average rest-frame colors of E+A galaxies located in the field and in clusters. Results show that average rest-frame near-IR colors of E+A galaxies located in clusters at $z \\sim 0.05$ \\citep{caldwell97} and field E+As located at $z \\sim 0.1$ (from the LCRS, Zabludoff \\etal 1996), follow the same color trend in $J - H$ and $H - K_s$ observed in the comoving color-color diagram: E+A galaxies located in nearby clusters appear bluer than field E+As ($z \\sim 0.1$).   As well as comparing the observed colors with the GISSEL96 colors at different redshifts, the models do not fit the rest-frame colors of the E+A galaxies observed in clusters at $z \\sim 0.3$. Their $H -K_s$ colors appear bluer (using the K-corrections of PEGASE), or redder (using the K-corrections of Poggianti) compared to the models. Their $J - H$ color index, although not particularly sensitive to one or the other K-correction, is also redder than the colors predicted by the models. The color of control galaxies, most of them ellipticals at $z \\lesssim 0.01$, and the others from clusters at $z \\sim 0.05$, agree with the near-IR colors predicted by models.  Integrated colors between the {\\em field} E+As and the {\\em cluster} E+As of the LCRS (LCRS \\# 4, 11 and 20; see Table \\ref{all_calib}) are similar, although those in clusters seem to be slightly ($\\sim 25\\%$) bluer in $H - K_s$ than the average color. This result is the same for both sets of rest-frame colors, the set corrected by the PEGASE K-corrections and the set corrected by the Poggianti K-corrections (see Figure \\ref{color_average}). On the other hand, the corresponding $J - H$ color is similar for the cluster and field E+As in subsample 1.  In order to build more robust results, more field and cluster E+A galaxies have to be observed  between $z = 0.1 - 1.0$. Spectroscopic observations of normal and E+A galaxies at different redshifts in the near-IR are necessary to (1) obtain calibrated SEDs and realistic K-corrections, and (2) to compare the spectra of the E+A galaxies with those of normal galaxies in the whole spectral range 3500 \\AA~ $\\lesssim \\lambda \\lesssim 25000$ \\AA. We expect to continue this research by imaging new E+A galaxies in the near-IR at higher redshift, as well as obtaining near-IR spectra in order to construct a useful and larger database of normal and post-starburst galaxies in a large spectral range. In order to increase the number of E+A galaxies, some field galaxies already classified as E+As, are being observed in the near-IR at Las Campanas. Some of these galaxies belong to the ESO-Sculptor Survey \\citep{delapparent97}, and results will be published soon. Other wide-field surveys will provide a wealth of data for E+A galaxies at $0.01 \\lesssim z \\lesssim 0.2$, like SLOAN \\citep{loveday98,fan98}, and the 2DF survey \\citep{colless98}, whose data are expected to become available to the public. In a forthcoming paper, we shall investigate systematic properties on the surface photometry and colors of the E+A galaxies."
    },
    "0002/astro-ph0002370_arXiv.txt": {
        "abstract": "We study the effect of a single, instantaneous starburst on the dynamical and chemical evolution of a gas-rich dwarf galaxy, whose potential well is dominated by a dark matter halo.  We follow the dynamical and chemical evolution of the ISM by means of an improved 2-D hydrodynamical code coupled with detailed chemical yields originating from type II SNe, type Ia SNe and single low and intermediate mass stars (IMS). In particular we follow the evolution of the abundances of H, He, C, N, O, Mg, Si and Fe.  We find that for a galaxy resembling IZw18, a galactic wind develops as a consequence of the starburst and it carries out of the galaxy mostly the metal-enriched gas. In addition, we find that different metals are lost differentially in the sense that the elements produced by type Ia SNe are more efficiently lost than others. As a consequence of that we predict larger [$\\alpha$/Fe] ratios for the gas inside the galaxy than for the gas leaving the galaxy. A comparison of our predicted abundances of C, N, O and Si in the case of a burst occurring in a primordial gas shows a very good agreement with the observed abundances in IZw18 as long as the burst has an age of $\\sim 31$ Myr and IMS produce some primary nitrogen.  However, we cannot exclude that a previous burst of star formation had occurred in IZw18 especially if the preenrichment produced by the older burst was lower than $Z=0.01$ Z$_{\\odot}$. Finally, at variance with previous studies, we find that most of the metals reside in the cold gas phase already after few Myr. This result is mainly due to the assumed low SNII heating efficiency, and justifies the generally adopted homogeneous and instantaneous mixing of gas in chemical evolution models. ",
        "introduction": "Dwarf irregular galaxies (DIG) are playing an increasingly central role in understanding galaxy evolution. This kind of galaxies generally has a low metallicity (from 0.5 Z$_{\\odot}$ to 0.02 Z$_{\\odot}$), a high gas content (up to $\\sim$ 10 times the stellar content) and their stellar populations appear to be mostly young. All these features indicate that these galaxies are poorly evolved objects.  Many gas-rich dwarf galaxies are known to be in a starburst phase, or are believed to have experienced periods of intense star formation in the recent past. These galaxies are generally called blue compact dwarf (BCD) galaxies (Sandage \\& Binggeli 1984). In general, dwarf gas-rich galaxies, given their simple structures and small sizes, are excellent laboratories to investigate the feedback of starbursts on the interstellar medium (ISM), and to study their chemical evolution. The aim is to reproduce the observed abundance ratios, to trace their recent star formation history and to discover if these galaxies could be the source of the intracluster gas (Gibson \\& Matteucci 1997).  Many authors have tried to connect late-type gas-rich (DIG and BCD) and early-type gas-poor dwarf galaxies (dwarf ellipticals and dwarf spheroidals) in an unified evolutionary scenario. The favourite theory about ISM depletion in gas-rich dwarf galaxies is based on the starburst-driven mass loss. The basis of this model, proposed by Larson (1974) and then applied specifically to dwarfs by Vader (1986) and Dekel \\& Silk (1986), is that the ISM is blown out of the galaxy by the energetic events associated with the star formation (stellar winds and supernovae). The well-known correlation between mass and metallicity found for both late-type and early-type dwarf galaxies (Skillman, Kennicut \\& Hodge 1989) is a natural result of the increasing inability of massive galaxies to retain the heavy elements produced in each stellar generation. At the present time is not yet clear if galactic winds are really the key point for understanding the formation and evolution of dwarf galaxies (see Skillman \\& Bender 1995 and Skillman 1997 for critical reviews about this point), but they certainly play an important role, regulating the mass, metal enrichment and energy balance of the ISM.  Observational evidences in support of the presence of outflows have been found recently in a lot of gas-rich dwarf galaxies, like NGC1705 (Meurer et al. 1992), NGC1569 (Israel 1988), IZw18 (Martin 1996) and many others. In their search for outflows in dwarf galaxies, Marlowe et al. (1995) pointed out that this kind of phenomena is relatively frequent in centrally star-forming galaxies. Again, they note a preferencial direction of propagation along the galaxy minor axis. In spite of these observational evidences, it is often difficult to estabilish if the gas will leave definitively the parent galaxy. In order to understand the final fate of both the swept-up gas and the metals ejected during the starburst and to study possible links between early and late-type dwarfs, numerical simulations are needed.  There are a lot of recent hydrodynamical simulations concerning the behaviour of the ISM and the metals ejected by massive stars after a starburst.  These simulations generally agree on the fact that galactic winds are not so effective in removing the ISM from dwarf galaxies, but disagree on the final fate of the metal-enriched gas ejected by massive stars. Many authors (D'Ercole \\& Brighenti 1999, hereafter DB; MacLow \\& Ferrara 1999, hereafter MF; De Young \\& Heckman 1994; De Young \\& Gallagher 1990) have found that galactic winds are able to eject most of the metal-enriched gas, preserving a significant fraction of the original ISM. Other authors (Silich \\& Tenorio-Tagle 1998; Tenorio-Tagle 1996) have suggested that the metal-rich material is hardly lost from the galaxies, since it is at first trapped in the extended haloes and then accreted back on to the galaxy.  However, all these models consider only the effects of stellar winds and SNII explosions on the dynamics of the ISM. In this paper we present models which take into account also the energetic contribution and the feedback from intermediate-mass stars and SNeIa, using the most up-to-date supernova rates. The effect of SNIa explosions is certainly fundamental for the late dynamical evolution of the ISM (up to $\\sim$ 500 Myr after the burst), even if their number is small.  There is an extensive literature about the chemical evolution of starburst and blue compact dwarf galaxies (see e.g. Matteucci \\& Chiosi 1983; Matteucci \\& Tosi 1985; Olofsson 1995). Pilyugin (1992, 1993) and Marconi et al. (1994) suggested the idea that the spread in the chemical properties of these galaxies, in particular the observed spread in He/H vs. O/H and N/O vs. O/H, could be due to self-pollution of H\\,{\\sc ii} regions coupled with `enriched' or `differential' galactic winds.  Therefore it is interesting to test the differential wind hypothesis with an hydrodynamical approach.  In our models we are able to follow the evolution in space and time of the abundances of several chemical elements (H, He, C, N, O, Mg, Si, Fe); in particular we follow, with suitable tracers, the gas released by stars of different initial mass. The chemical composition of each of these tracers is obtained by adopting the nucleosynthesis prescriptions from various authors (Woosley \\& Weaver 1995, hereafter WW; Renzini \\& Voli 1981, hereafter RV; Nomoto, Thielemann \\& Yokoi 1984, hereafter NTY).  In section 2 we describe the model and the assumptions adopted in our simulations. The results are presented in section 3 and compared with the observational constraints available for the BCG IZw18. A discussion is presented in section 4 while some conclusions and future improvements of the model are discussed in section 5. ",
        "conclusions": "We have studied the dynamical and chemical evolution of a dwarf galaxy as due to the effect of a single, instantaneous, point-like starburst occurring in its centre. We adopted galactic structural parameters which resemble those of IZw18, the most unevolved dwarf blue compact galaxy known locally.  We ran different models, which differ for the burst luminosity and the ISM mass. We considered the mass and energy inputs from the single low and IMS, the SNeII and the SNeIa (white dwarfs in binary systems) and we followed the evolution of the gas and its chemical abundances (H, He, C, N, O, Mg, Si and Fe) in space and time for several hundreds Myr from the burst.  Our results can be summarized as follows:  \\begin{enumerate}  \\item The starburst can inject enough energy into the ISM to trigger a metal enriched galactic wind: the metals synthetized and ejected through supernova explosions leave the galaxy more easily than the unprocessed gas.  This result is not new since it has already been suggested by previous works (e. g. MF and DB).  However, our new result is that the SNIa products have the largest ejection efficiency (more than the products of type II SNe), with the consequence that the [$\\alpha$/Fe] ratios in the gas outside the galaxy are predicted to be lower than those inside. This is due to the fact that SNeIa produce a substantial fraction of iron.  \\item The energy injection in the ISM by SNII has a rather low efficiency. Instead, the energetic contribution of SNeIa, in spite of their relatively small number (a total of 240 SNeIa against 4800 SNeII), has important consequences in the dynamical behaviour of the galaxy. Since the SNIa explosions occur in a medium heated and diluted by the previous activity of SNeII, the thermal energy of the explosions is easily converted into kinetic energy and so the gas reaches quickly the outer regions of the bubble along the galactic chimney. DB showed that after the end of SNII activity, a fraction of the gas tends to recollapse toward the central region of the galaxy, achieving the threshold density for a new star formation event ($N({\\rm H}\\,{\\sc i})\\simgt 10^{21}\\;{\\rm cm}^{-2}$, Skillman et al. 1988; Sait\\=o et al. 1992) in $\\sim$ 0.5--1 Gyr. With the energetic contribution of SNeIa, it is not possible to reach, at least for the time considered in our simulation, such a threshold in column density (it can be obtained only after many Gyr).  \\item One single burst, occurring in a primordial gas, with an age of $\\sim$ 31 Myr reproduces quite well both the dynamical structures and the abundances in IZw18. This value is consistent with other independent age estimates for the burst (Martin 1996).  From the nucleosynthetic point of view the age of 31 Myr ensures that there is enough time for the primary N from IMS to be produced and ejected. The adopted yields for massive stars (Woosley \\& Weaver 1995) include some primary N, but its amount is negligible. This result suggests that IZw18 is probably experiencing its first major burst of star formation, although we cannot exclude a previous burst (see Aloisi et al. 1999) of moderate intensity, which enriched the gas to a metallicity $Z < 0.01$ Z$_{\\odot}$.  \\item At variance with previous studies we find that the majority of metals (in mass) are found in the cold gas. In fact, mainly because of the low SNII efficiency, the wind superbubble remains several hundreds of Myr inside the galaxy before the break out occurs. Moreover, the superbubble becomes radiative after a few Myr, and most of the SNII ejecta cool without leaving the galaxy (SNIa ejecta, instead, are shown to be vented more easily). Given the relatively short mixing time, the abundances predicted by our models for the cold gas are those that should be compared with the abundances observed in IZw18.  Actually, they are in very good agreement with the observed ones.  This result supports the common assumption made in chemical evolution models of an instantaneous mixing.   \\end{enumerate}  Future improvements of this work will include models with continuous bursts, sequential bursts, as well as a more detailed study of the formation of the H\\,{\\sc ii} regions (Recchi et al. in preparation)."
    },
    "0002/astro-ph0002437_arXiv.txt": {
        "abstract": "We present the results of weak gravitational lensing statistics in four different cosmological $N$-body simulations. The data has been generated using an algorithm for the three-dimensional shear, which makes use of a variable softening facility for the $N$-body particle masses, and enables a physical interpretation for the large-scale structure to be made. Working in three-dimensions also allows the correct use of the appropriate angular diameter distances.  Our results are presented on the basis of the filled beam approximation in view of the variable particle softening scheme in our algorithm. The importance of the smoothness of matter in the universe for the weak lensing results is discussed in some detail.  The low density cosmology with a cosmological constant appears to give the broadest distributions for all the statistics computed for sources at high redshifts. In particular, the range in magnification values for this cosmology has implications for the determination of the cosmological parameters from high-redshift Type Ia Supernov\\ae. The possibility of determining the density parameter from the non-Gaussianity in the probability distribution for the convergence is discussed. ",
        "introduction": "\\subsection{Outline}  We present the results of a study of the weak gravitational lensing of light in four different cosmological models, using the algorithm for the three-dimensional shear developed by Couchman, Barber and Thomas (1999). Since weak lensing effects depend on the angular diameter distances for the lenses and sources, and also the specific distribution and evolution of matter, the results are sensitive to the particular cosmological model.  In {\\em strong} lensing studies, frequent use is made of the `thin-screen approximation,' in which the mass distribution of the lens is projected along the line of sight and replaced by a mass sheet with the appropriate surface density profile. Deflections of the light from the source are then considered to take place only within the plane of the mass sheet, making computations for the light deflections simple.  The simplicity of the thin-screen approximation has also lead to its frequent use in weak gravitational lensing studies, where each of the output volumes from cosmological $N$-body simulations is treated as a planar projection of the particle distribution within it.  However, a number of problems can arise with two-dimensional approaches, especially in weak lensing studies in which the large-scale distribution of matter extending to high redshifts is responsible for the lensing. Couchman, Barber and Thomas (1999) considered some of the shortcomings and were motivated to develop an algorithm to compute the six independent components in three dimensions of the second derivative of the gravitational potential (the three-dimensional shear). In the method used in the present work, for which the effects of lensing along lines of sight are required, the two-dimensional `effective lensing potentials' (see Section 2) are obtained by integrating the computed three-dimensional shear components along the lines of sight. These effective lensing potentials are used to construct the Jacobian matrices which are recursively generated along the lines of sight, and from the final Jacobian matrices the magnifications and two-dimensional shear are determined.  A brief outline of this paper is as follows.  In {\\bf Section 1.2} we summarise the previous weak gravitational lensing methods by other authors.  In {\\bf Section 2}, the essential equations for gravitational lensing are explained, and the multiple lens-plane theory is described, culminating in expressions for the magnification, convergence, two-dimensional shear, and ellipticity, which are the weak lensing outputs from the cosmological simulations.  {\\bf Section 3.1} summarises the three-dimensional shear algorithm, and particularly the choice of the variable softening scale for the particles. This feature limits the effects of isolated particles (not representative of the large-scale structure), and builds in a physical significance both to the choice of the softening scale, and to the underlying form of the dark matter. We explain also how the softening carefully limits the incidence of strong lensing events in all the cosmologies. In {\\bf Section 3.2} we describe the four cosmological $N$-body simulations used in this work. In {\\bf Section 3.3} we state how the simulation boxes are oriented to minimise correlations in the large-scale structure between adjacent time-outputs, and we describe the establishment of lines of sight through the simulations, and the locations of the evaluation positions for the shear within them. In {\\bf Section 3.4} we explain how the computed shear values are converted (after integration along the lines of sight) to physical units, and the approximations we have adopted in the method.  In {\\bf Section 4.1} the Dyer-Roeder equation for the angular diameter distances is introduced and {\\bf Section 4.2} summarises the results derived in the Appendix for the beam equation generalised for all the cosmologies. This precedes the discussion ({\\bf Section 4.3}) of magnifications obtainable in inhomogeneous universes with different degrees of smoothness, and highlights the significance of differences between researchers who adopt different approaches to this subject.  {\\bf Section 5} formally presents the weak lensing results. We first ({\\bf Section 5.1}) attempt to see if the onset of structure formation can be seen from the shear data, and briefly comment on the expected behaviour of the developing shear. Secondly ({\\bf Section 5.2}), the results for magnification, convergence, shear and ellipticity are presented for all the cosmologies, compared and contrasted. We consider the impact of the smoothness parameter on the results.  It may be possible to determine the density parameter, $\\Omega_0$, from the probability distribution and the skewness in the distribution for the convergence, both of which may be measurable observationally. In {\\bf Section 6.1} we present our results for these quantities, and compare them with others. The application to the determination of the cosmological parameters from Type Ia Supernov\\ae~ is mentioned in {\\bf Section 6.2}.  In {\\bf Section 7} the weak lensing statistics detailed in Section 5 are summarised, together with the results of the non-Gaussianity in the convergence. We also compare and contrast our results with other authors. Of particular significance are the common use of two-dimensional approaches by others, and the use of either point mass particles or particles with small softening scales, which introduce high values of magnification. These authors often use the empty cone approximation, rather than the full beam approximation, making comparisons difficult. This latter point is discussed.  In the {\\bf Appendix} we state the generalised beam equation, and derive the equations necessary for the numerical determination of the angular diameter distances for all the cosmologies.  \\subsection{Previous work}  Numerous methods have been employed to study weak lensing, and throughout this paper we will make comparisons with previous work by other authors. We summarise here the methods which have been used by others.  Jaroszy\\'{n}ski et al. (1990) use a `ray-tracing' method to evaluate the matter column density in a matrix of 128$^3$ pixels for each of their lens-planes. The boxes were generated using a particle-mesh (PM) code in the standard Cold Dark Matter (SCDM) cosmology, and were of side-dimension 128$h^{-1}$Mpc, where $h$ is the Hubble constant in units of 100~km~s$^{-1}$~Mpc$^{-1}$. By making use of the assumed periodicity in the particle distribution orthogonal to the line of sight, they translate the planes for each ray, so that it becomes centralised within the plane of one full period in extent. This removes any bias acting on the ray when the shear is computed. Instead of calculating the effect of every particle on the rays, the pixel column densities in the single period plane are used, and they assume that the matter in each of the pixels resides at the centre point of each pixel. They calculate the two two-dimensional components of the shear (see Section 2 for the definition of shear) as ratios of the mean convergence of the beam, which they obtain from the mean column density. However, they have not employed the net zero mean density requirement in the planes, (described in detail by Couchman, Barber and Thomas, 1999), which ensures that deflections and shear can only occur when there are departures from homogeneity. Also, the matter in the pixel through which the ray is located is excluded. To follow the shearing across subsequent planes they recursively generate the developing Jacobian matrix for each ray, in accordance with the multiple lens-plane theory (see Section 2).  Wambsganss, Cen and Ostriker (1998) also use the `ray-tracing' method in cosmological $N$-body simulations. The method is applied to PM simulations of the SCDM cosmology, in which a convolution method is used to combine large-scale boxes of 400$h^{-1}$Mpc and resolution 0.8$h^{-1}$Mpc, with small-scale boxes of 5$h^{-1}$Mpc and the higher resolution of 10$h^{-1}$kpc. They randomly orient each simulation box, and project the matter contained within each onto a plane divided into pixels. They choose the central $8h^{-1}$Mpc $\\times~8h^{-1}$Mpc region through which to shoot rays, but account for the deflections of the rays in terms of all the matter in the plane of $80h^{-1}$Mpc $\\times ~80h^{-1}$Mpc. However, to speed up the computation, a hierarchical tree code in two dimensions is used to collect together those lenses (pixels) far away, whilst treating nearby lenses individually. The code assumes that all the matter in a pixel is located at its centre of mass. The matter in each pixel, which measures $10h^{-1}$kpc $\\times~10h^{-1}$kpc, is assumed to be uniformly spread.  By using the multiple lens-plane theory, they show both the differential magnification probability distribution, and the integrated one for 100 different source positions at redshift $z_s=3.0$. One advantage of this type of ray-tracing procedure is its ability to indicate the possibility of multiple imaging, where different rays in the image plane can be traced back to the same pixel in the source plane, and they are able to compute the statistics of angular separations for multiple images.  Marri and Ferrara (1998) use lens-planes up to redshifts of $z=10$ for mass distributions determined by the Press-Schechter formalism, and treat the particles as point-like masses with no softening. They apply their ray-tracing method to three cosmologies with $(\\Omega_M,~\\Omega_{\\lambda},~\\Omega_{\\nu}) = (1,~0,~0)$, (0.4, 0.6, 0), and (0.7, 0, 0.3). ($\\Omega_M$, $\\Omega_{\\lambda}$ and $\\Omega_{\\nu}$ represent the density parameters for matter, vacuum energy and the hot Dark Matter component, respectively.)  The maximum number of lenses in a single plane is approximately 600, each having the appropriate computed mass value, and they follow $1.85 \\times 10^7$ uniformly distributed rays within a solid angle of $2.8 \\times 10^{-6}$sr, corresponding to a $420'' \\times 420''$ field. The final impact parameters of the rays are collected in an orthogonal grid of $300^2$ pixels in the source plane. Because of the use of point masses, their method produces very high magnification values, greater than 30 for the Einstein-de Sitter cosmology. They have also chosen to use a smoothness parameter $\\bar{\\alpha} =0$ in the redshift-angular diameter distance relation (described in Section 4) which depicts an entirely clumpy universe.  Jain, Seljak and White (1998, 1999) have made use of $N$-body simulations generated using a parallel adaptive particle-particle, particle-mesh (AP$^3$M) code. The cosmologies simulated are the same as the ones reported on here, but they have 256$^3$ particles and comoving box side dimensions of 85$h^{-1}$Mpc for their Einstein-de Sitter cosmologies, and 141$h^{-1}$Mpc for the open and cosmological constant models. For each plane, the projected density, together with the appropriate redshift-dependent factors, is Fourier transformed, using periodic boundary conditions, to obtain the shear in Fourier space. The two-dimensional shear matrix is then computed onto a grid in real space by the inverse Fourier transform.  The size of the grid is chosen to be compatible with the force-softening scale (the resolution) in the $N$-body simulations, so that for sources at redshifts around 1, the angular-scale size of the grid is less than, or of the same order as, the angular-scale size of the force-softening at these redshifts. Perturbations on the photon trajectories are computed and the shear matrix interpolated to the photon positions enabling the Jacobian matrices to be computed by recursion. In view of the use of a fine grid for the shear and deflection angle computations, they are able to analyse their data on different scales, and are thus able to determine the power spectrum in both the shear and the convergence. This approach is very different from the approach we follow which makes use of variable softening for the particles depending on their particular environment, and which ensures that most `rays' would pass entirely through regions of smoothed density, thereby requiring the application of the full beam approximation. Jain, Seljak and White (1998, 1999) assume point particles interpolated onto their grid, and use the empty cone approximation.  Hamana, Martel and Futamase (1999) study weak lensing in $N$-body simulations of three cosmologies with ($\\Omega_0,~\\lambda_0$) = (1, 0), (0.3, 0.7), and (0.3, 0). (We now use $\\Omega_0$ and $\\lambda_0$ to represent the present matter density and vacuum energy density parameters respectively.) These were generated by a P$^3$M algorithm, using Fourier techniques on a 128$^3$ lattice. The comoving simulation box sizes were 128Mpc, and the particles were given comoving softenings of 300kpc. Three simulations were performed for each cosmological model, and boxes from the different simulations combined to limit correlations in large-scale structure between adjacent boxes. Each of the mass distributions in the simulation boxes were than projected onto planes at the box redshift, and Poisson's equation solved numerically on each of them. This was done by first evaluating the surface density onto a $512~\\times~512$ grid, based on the particle positions, and then inverting Poisson's equation using a Fast Fourier Transform (FFT) method. As the multiple lens-plane theory was to be used for more than $10^7$ rays passing through each plane, the recursion algorithm was simplified by assuming small deflections for the rays. This also meant that the rays could be considered to pass through the grid points, and travel, effectively, in straight lines through the entire distance from the observer to the source plane. We have also made this approximation in our own method.  An alternative to the conventional form of `ray-tracing' was introduced by Refsdal (1970), who used `ray-tracing' with calculations of the differential deflections of light rays around a central ray to determine the distribution of magnifications.  Fluke, Webster and Mortlock (1999) and Fluke, Webster and Mortlock (2000) have developed this idea further as the `ray-bundle' method. The principle is to trace the passage of a discrete bundle of light rays as it passes through the deflection planes. The advantage of the method is that it provides a direct comparison between the shape and size of the bundle at the observer and at the source plane, so that the magnification, ellipticity and rotation can be determined straightforwardly. The authors have selected a range of popular cosmologies, including the ones reported on here, and have produced their own $N$-body simulation data-sets using a P$^3$M algorithm. For each cosmology they have produced a number of independent simulations to enable them to randomly choose a simulation box for a particular epoch from any of the realisations. In this way, correlations of large-scale structure between adjacent boxes are avoided. They have run their simulations with $64^3$ dark matter particles, with box sizes ranging from 80$h^{-1}$Mpc to 164.3$h^{-1}$Mpc. Having randomly selected the boxes for a given cosmology, the particle mass distributions are then projected onto planes at the redshifts of the boxes. The authors construct approximately 50,000 bundles, each comprising 8 rays, to shoot in random directions through the planes, from the observer's location at $z = 0$. The shooting area was limited to $50'' \\times 50''$ to avoid edge effects of the planes, and only matter within a single period in the transverse direction is included in the planes. In addition, a radius, typically 15$h^{-1}$Mpc centred around each bundle is chosen for the extent of the matter to be included in the deflection angle computations. The projected masses are also considered as point particles, so that very high magnification values can be achieved in principle; however, the authors do not include bundles which pass within $\\sqrt{2}$ of the Einstein radius of any particle. Because of their use of point masses, they use the empty cone approximation in the determination of the statistical distributions of magnifications for the different cosmological models.  Premadi, Martel and Matzner (1998a) have introduced individual galaxies into the computational volume, matching the 2-point correlation function for galaxies. They also assign morphological types to the galaxies according to the individual environment, and apply a particular surface density profile for each. Five different sets of initial conditions were used for the simulations, so that the individual plane projections can be selected at random from any set. $N$-body simulations were produced for three cosmologies with ($\\Omega_0,\\lambda_0$) = (1, 0), (0.2, 0), and (0.2, 0.8) in boxes of comoving side-dimension 128$h^{-1}$Mpc. They solve the two-dimensional Poisson equation on a grid and invert the equation using a FFT method to obtain the first and second derivatives of the gravitational potential on each plane. They also correctly ensure that the mean surface density in each lens-plane vanishes, so that a good interpretation of the effects of the background matter is made. Their method uses beams of light, each comprising 65 rays arranged in two concentric rings of 32 rays each, plus a central ray. The multiple lens-plane theory then enables the distributions of cumulative magnifications to be obtained.  Tomita (1998a and b) also uses a ray bundle method, but by evaluating the gravitational potential at approximately 3000 locations between the observer and sources at redshift 5, he is able to compute the weak lensing statistics without using the multiple lens-plane theory. He has used $N$-body simulations produced using a tree code with $32^3$ particles in four different cosmologies with $(\\Omega_0,~\\lambda_0) = (1,~0),~(0.2,~0.8),~(0.2,~0),$ and (0.4, 0). In all of these, except the SCDM cosmology ($\\Omega_0 = 1,~\\lambda_0 = 0$), the particles are treated as galaxy-size objects. In the SCDM cosmology, 20\\% of the particles are treated thus, whilst 80\\% are given softening radii of 100$h^{-1}$kpc (model a), 40$h^{-1}$kpc (model b), and 20$h^{-1}$kpc (model c). Ray bundles are formed from $5 \\times 5$ rays arranged in a square, with separation angles (fixed for each run) at between 2 arcsec and 1 degree. At each of the 3000 positions between the observer and source, the potential is evaluated at each ray position by translating each simulation cube (using the periodic properties) so that the centre of each bundle is located in the centre of each cube. In this way, all the matter within a full period in extent contributes to the calculation of the potential (although no account is taken of matter beyond one period transverse to the line of sight). To avoid spurious values of the potential arising from masses close to evaluation positions, an average of the potential is taken by integrating it analytically over the interval between adjacent evaluation positions. The light propagation is then determined by solving the null-geodesic equations, and the required statistics constructed from shooting 1000 bundles through the flat cosmologies, and 200 bundles through the open cosmologies. ",
        "conclusions": "In the application of the code for the three-dimensional shear, we have had to consider what appropriate angular diameter distance values should be applied to the data. The values have all been calculated numerically from the generalised beam equations (see section 4) for the different cosmologies. This was done in each case for a smoothness parameter $\\bar{\\alpha} = 1$. Our variable softening scheme for the particles ensures that nearly all rays pass entirely through softened mass, and in addition, we found that the minimum value of $\\bar{\\alpha}$ was at least 0.8 (at $z = 0$) in all the cosmologies. The differences in the magnification distributions for $\\bar{\\alpha} = 1$ and for the minimum value were almost indistinguishable. The effects of shear on the angular diameter distances (through changes to the effective value of the smoothness parameter, $\\bar\\alpha$), were found to be completely negligible at all redshifts, so we are justified in ignoring them in the distance-redshift relation. Furthermore, they are always additive, making the effective value of $\\bar{\\alpha}$ even closer to unity.  We found, in Section 5.1, that the Einstein-de Sitter universes showed the most rapid growth in the `intrinsic' shear components at late times, as expected from the growth of structure in these cosmologies. This was as a result of studying the computed shear values before the application of the angular diameter distance factors, and before conversion to physical units.  The OCDM and LCDM cosmologies indicated only a limited contribution from the cosmological constant in terms of the growth in the shear values, again consistent with the expected evolution of structure in these cosmologies.  When the computed shear values were multiplied by the full conversion factors appropriate to the integration, together with the angular diameter distance factors, the resulting curves exhibited very broad peaks, indicating that significant lensing may result from structure in a wide band of redshifts. Significantly, the LCDM cosmology has both the broadest and the highest peak, indicating that this cosmology should produce, for example, the broadest range of magnifications. This result would appear to come primarily from the large values of the angular diameter distance factors for this cosmology, rather than any considerations about the evolution of structure. However, the broader and higher peak for the SCDM cosmology, compared with the TCDM cosmology, does indicate the differences in structure within them, since they both have the same values for the angular diameter multiplying factor, $R$.  In Section 5.2, we showed the results for the magnification distributions for the different cosmologies for different source redshifts, and these are concisely summarised in Table~\\ref{mu4C}. At high redshift, the LCDM cosmology produces the highest magnifications, the broadest distribution curves, and the lowest peak values. For sources at $z_s=3.6$ in the LCDM cosmology, 97\\mbox{$\\frac{1}{2}$}\\% of all lines of sight have magnification values up to 1.850. (The maximum magnifications, not quoted here, depend on the choice of the minimum softening in the code, although the overall distributions are very insensitive to the softening.) The rms fluctuations in the magnification (about the mean) were as much as 0.191 in this cosmology, for sources at $z_s=3.6$.  Even for sources at $z_s=0.5$ there is a measurable range of magnifications in all the cosmologies.  The immediate implication of these results is the likely existence of a bias in observed magnitudes of distant objects, and a likely dispersion for standard candles (for example, Type Ia Supernov\\ae) at high redshift.  The magnification versus the convergence showed the presence of significant shear, and the mean values of $\\gamma$ in small convergence bins pointed to a possible slow linear increase in $\\langle \\gamma \\rangle$ with $\\kappa$. This would result in a trend towards higher mean magnification values (as $\\kappa$ increases) than would be the case for constant $\\langle \\gamma \\rangle$.  The distributions for the shear, $\\gamma$, are broadest, in the LCDM cosmology for high source redshifts, and there are closely linear relationships between the ellipticity and the shear in all the cosmologies. This relationship leads again to broad distributions in the ellipticity for high-redshift sources in the LCDM cosmology. The peak in the ellipticity distribution for the LCDM cosmology, for $z_s = 3.6$, is 0.111, being almost twice the value in the TCDM cosmology.  Jain, Seljak and White (1999) have expressed the possibility of determining the value of the density parameter from the convergence field from weak lensing statistics. They show one-point distribution functions for $\\kappa$, assuming sources at $z_s = 1$, for all their four cosmologies and using different (fixed) smoothing scales in each. They describe the increasing non-Gaussianity of the distribution functions as the smoothing scale is reduced, and the increasing tail at high $\\kappa$. They also describe the shape of the distribution functions for negative $\\kappa$, which results from the rate of structure formation in the different cosmologies, and claim the interesting conclusion that the minimum value of $\\kappa$ is proportional to the density parameter. We were able to establish from our own work that, in broad terms, the skewness in $\\kappa$ decreases with source redshift, and decreases with increasing density parameter, as expected. Also, the statistic $S_3$ (defined in equation~\\ref{S3}) was found to decrease with increasing $\\Omega_0$. However, we were unable to establish more precise relationships, because of the use of the full beam approximation. The order of the cosmologies for $S_3$ is the same as that presented by Jain, Seljak and White (1999). We also found that both $\\kappa_{\\mathrm{min}}$ and $(\\langle \\kappa \\rangle -\\kappa_{\\mathrm{min}})$ are in the correct order for the different cosmologies in terms of $\\Omega_0$, although, again for the same reason, it would not be possible to obtain specific values for $\\Omega_0$.  We now, very briefly, make some comparisons with the weak lensing results obtained by other authors. This will not be exhaustive, because it has been anticipated that our results should be different, for a number of reasons. Primary amongst these are the following. First, our results were obtained using the three-dimensional shear code, which allows periodicity, the use of the peculiar potential, the net zero mean density requirement, and angular diameter distances to every evaluation position within each simulation volume. Second, many two-dimensional (planar) approaches may suffer from inadequate convergence to the true limiting values for the shear matrix and angular deflections. Third, we have introduced a physically realistic variable softening to the method, which requires use of the full beam approximation for the angular diameter distances, rather than the empty cone approximation used by many authors with either point masses, small (fixed) softenings, or small pixellation in the planes.  The magnification distributions of Jaroszy\\'nski et al. (1990) for the SCDM cosmology do not have mean values of 1, and their dispersions in the convergence for sources at $z_s=1$ and $z_s=3$ are considerably lower than our values, with very little evolution with redshift.  In our work, the peaks of the ellipticity distributions occured at values of 0.075 (similar to the value for $z_s=4$) and 0.034 for sources at $z_s =3$ and 1 respectively, and these are somewhat lower than the values of 0.095 ($z_s=3$) and 0.045 ($z_s=1$) found by Jaroszy\\'nski et al. (1990). Rather surprisingly, however, their peak values in the distributions for the shear are quite similar to ours, especially for sources at $z_s=3$.  Wambsganss, Cen and Ostriker (1998) find magnifications up to 100 in the SCDM cosmology, and correspondingly highly dispersed distributions, very much larger than ours for $z_s=3$.  The high magnification tail in the distributions almost certainly derives from the low value of the (fixed) softening scale resulting from the `smearing' of the mass distribution in the $10h^{-1}$kpc $\\times 10h^{-1}$kpc pixels.   Similarly, Marri and Ferrara (1998) show wide magnification distributions, and very high maximum values, again which occur as a result of using point particles rather than smoothed particles. Their procedure (summarised in Section 1.2) is completely different from ours, and, unlike us, they find that the SCDM cosmology has the broadest magnification probability distribution, followed by the LCDM cosmology, and finally their HCDM cosmology. In particular, we would disagree with their choice of $\\bar\\alpha = 0$, which is representative of an entirely clumpy universe, as opposed to our finding, that the SCDM universe is close to being smooth at all epochs.  Hamana, Martel and Futamase (1999) find that the dispersions in the probability distributions for $\\kappa$, $\\gamma$ and $\\mu$ are all greatest for the Einstein-de Sitter cosmology, and are very similar for their open and cosmological constant cosmologies. This is in complete contrast to our findings. Their magnification and convergence distributions for $z_s = 1,$ 2 and 3 are much broader than ours, although the distributions for $\\gamma$ in the SCDM cosmology are similar.  The ranges in magnification from Premadi, Martel and Matzner (1998a) appear to be rather similar to ours for sources at $z_s=3$ for their three cosmologies, and the widths of the distributions are in the same order for the different cosmologies that we find. This is reassuring because, although their method relies on two-dimensional projections of the simulation boxes, they include many of the essential features to which we have drawn attention, for example, an assumed periodicity in the matter distribution, randomly chosen initial conditions to avoid structure correlations between adjacent simulation boxes, the net zero mean density requirement, realistic mass profiles for the particles, and use of the filled beam approximation with a smoothness parameter, $\\bar\\alpha =1$. They show the average shear for a source at $z_s=5$ contributed by each of the lens-planes individually, and find that the largest contributions come from those planes at intermediate redshift, of order $z$=1 -- 2. Similarly, they find that the lens-planes which contribute most to the average magnifications are also located at intermediate redshifts. In terms of the development of structure in the different cosmologies, they find that the lens-planes contributing the most shear and magnification are located at larger redshifts for those cosmologies with smaller $\\Omega_M$. The average shear for each redshift has been plotted by the authors. They find, for all redshifts, that the values in the SCDM cosmology are many times greater than in the LCDM cosmology, which is, in turn, greater than in the OCDM cosmology. This is not what we find.  Fluke, Webster and Mortlock (2000), in applying the ray bundle method of Fluke, Webster and Mortlock (1999), explain clearly the differences between the empty cone and full beam approximations. They use the empty cone approximation, because of their use of an effective fixed physical radius for each particle, equal to $\\sqrt{2}\\times$ the Einstein radius for each. This gives rise to magnification probability distributions with $\\mu_{\\mathrm{min}} = 1$, arising from the use of $\\bar{\\alpha}=0$, and high magnification tails, arising from the small effective radii for particles and clusters. They obtain the weak lensing statistics for the same cosmologies we have used. For $z_s = 1$, the distribution in the magnifications for the SCDM cosmology is clearly broader than that for the LCDM cosmology, which in turn is broader than the OCDM cosmology. This is also true of our data, although the order for the cosmologies, in our work, is completely altered for higher source redshifts. The authors find, as do we, that the most significant differences amongst the cosmologies arise directly as a result of the optical depth to the source, which is related to the angular diameter distances. Moreover, there is a suggestion in their results (Fluke, 1999) that Kolmogorov-Smirnov tests may be unable to distinguish the cosmologies, {\\em even random distributions of particles,} in terms of their magnification distributions, if the same set of angular diameter distance factors is applied to each.  This last statement may prove to be extremely important, as it appears from both our own work and that of Fluke, Webster and Mortlock (2000), that the angular diameter distances (or optical depths) are really the determining factor for weak lensing statistics. Moreover, our two approaches may represent `limits' for the true lensing behaviour from the mass distribution in the universe. The variable softening facility within our algorithm leads naturally to the assumption that the universe may be described in terms of the full beam approximation. This is, however, quite different from the assumptions of most other workers, who frequently use point particles, or a limited form of fixed softening, or small pixellation, and therefore use the empty beam approximation. The two approaches give rise to quite different expectations and results. The most obvious differences are the following. First, strong lensing can occur with effectively small particles, leading to high magnification tails in the probability distributions. Second, magnification distributions in the empty cone approximation all have $\\mu_{\\mathrm {min}} = 1$, whilst in the full beam approximation, $\\mu_{\\mathrm {min}} \\leq 1$; this may alter the dispersions in the two distributions. Third, the mean values for the magnifications can be calculated from the respective angular diameter distances in the different cosmologies for the empty beam approximation; however, the mean values in the full beam approximation are always 1. These points make comparisons between methods using the different approximations difficult. However, it is probable that the universe is neither completely smooth nor filled with galactic-mass point-like objects. (Subramanian, Cen and Ostriker, 1999, highlight our uncertainty in this area, by suggesting that small dense masses formed early during hierarchical clustering may persist to late times, so that real cluster halo structures may depend crucially on the detailed dynamics of the dense pockets.) The resolution of this question together with a better understanding of the form, distribution and evolution of the dark and luminous matter content of the universe should provide a much clearer indication of the likely weak lensing statistics in cosmological $N$-body simulations of the future."
    },
    "0002/astro-ph0002092_arXiv.txt": {
        "abstract": "We simulate fragmentation and gravitational collapse of cold, magnetized molecular clouds.  We explore the nonlinear development of an instability mediated by ambipolar diffusion, in which the collapse rate is intermediate to fast gravitational collapse and slow quasistatic collapse. Initially uniform stable clouds fragment into elongated clumps with masses largely determined by the cloud temperature, but substantially larger than the thermal Jeans mass. The clumps are asymmetric, with significant rotation and vorticity, and lose magnetic flux as they collapse. The clump shapes, intermediate collapse rates, and infall profiles may help explain observations not easily fit by contemporary slow or rapid collapse models. ",
        "introduction": "The interstellar magnetic field plays an important role in the dynamics of molecular clouds and the collapse of dense cloud cores into protostars. Magnetic pressure and tension combine with thermal and turbulent kinetic pressure to resist gravitational collapse.  The role of the magnetic field is often simply characterized by a critical mass $M_{\\rm crit}$ which depends on the magnetic flux threading the cloud \\citep{ms56,ms76,tin88,mckee93,m99}. Clouds with $M > M_{\\rm crit}$ are termed supercritical and collapse on a dynamical time-scale: the magnetic field can have a moderate effect on the morphology of collapse and can slow collapse in the cloud envelope \\citep[{\\it e.g.} ][]{black82}, but cannot significantly slow collapse in the core. Clouds with $M < M_{\\rm crit}$ are termed subcritical, and evolve on a longer time-scale as magnetic support is lost due to ambipolar diffusion. The magnetic field is redistributed within the cloud so that the inner parts become supercritical. The cloud is then differentiated into a dynamically collapsing core with a magnetically supported envelope \\citep{cm93,fm93,bm94,cm94,sms97,ck98}. Much progress has been made in following this type of evolution through 6 or more orders of magnitude of increase in central density, including the effects of rotation as well as detailed chemistry and grain physics. The outcomes of these calculations include detailed density, bulk velocity, ion-neutral drift velocity, magnetic field, and grain and ion abundance profiles in axisymmetric clouds, as well as an appreciation of the timescales, rate of magnetic flux loss, and role of magnetic braking in this mode of isolated star formation.  This theoretical picture can be observationally tested \\citep[see recent reviews by ][]{evans99,meo99}. So far, the results are ambiguous. Flow toward an isolated infrared source in the Bok globule B335 is well described by an inside-out collapse model \\citep{z90,z93,z95}. The density distribution and measured magnetic fieldstrength in the cloud B1 have been fit by a model with a subcritical envelope and a core which has evolved to a supercritical state by ambipolar drift \\citep{c94}. On the other hand, in some respects the existing models appear to be incomplete. Clumps and cores are not axisymmetric; \\citet{m91} surveyed 16 dense cores a few tenths of a parsec in size and pointed out that at least 6 of them are likely to be prolate. \\citet{r96} made a statistical argument, based on a larger sample, that clumps and globules are more likely prolate or triaxial than oblate. \\citet{wt99} showed that asphericity appears also at smaller scales. Collapse guided by a magnetic field could produce oblate clouds, but not prolate ones. It also appears unlikely that rotation accounts for the flattening; this has been shown quantitatively in the case of L1527 \\citep{o97}. Thus, the shapes are unexplained. Moreover, in some cases in which infall has been measured directly, it is more spatially extended, with faster velocities in the outer parts, than expected from the standard models of inside-out collapse or gravitational motion driven by ambipolar drift. This has been shown in the case of L1544 by \\citet{t98}, and for 6 other starless cloud cores by \\citet{greg}. (Interestingly, a model of L1544 has been recently constructed by \\citet{cb2000} to match the observations of \\citet{t98}.  The model requires a mass-to-flux ratio more nearly critical than previously published models by the same authors, and appears to display the same intermediate collapse discussed in this paper, although the authors do not call it out as such.) Finally, there are observations which relate to the timescale for collapse of molecular clouds into protostars.  \\citet{lm99} find collapse timescales of $\\sim$ 0.3-1.6 Myr from the ratio of the numbers of starless cores to cores with embedded young stellar objects.  They state that this requires collapse 2-44 times faster than ambipolar drift models. These observations suggest that another ingredient may be required to explain the collapse of molecular clouds: the decay of turbulent support \\citep{mlazar}, or, as we explore in this paper, a magnetogravitational instability mediated by ambipolar drift.  The magnetic field likely also plays an important role in other aspects of collapse, which are still incompletely understood. The mechanism by which molecular clouds fragment, and the masses and morphology of those fragments, is clearly of importance to the stellar initial mass function and the origin of binary systems.  The magnetic field can exert strong forces on many scales, affecting fragmentation \\citep[{\\it e.g.} ][]{boss97,boss99}. The field may also be a source of kinetic energy in the cloud, if the free energy of an ordered field can be released as turbulent motions. These issues have not been studied thoroughly in the detailed axisymmetric  models of gravitational contraction with ambipolar drift referenced above because they would require calculations with no constraints on spatial symmetry.  In this paper, we address some of these issues by studying a simple problem: the evolution of small perturbations to an initially uniform, magnetically subcritical sheet of weakly ionized gas with a uniform magnetic field perpendicular to its plane. The perturbations evolve under the influence of magnetic tension, self gravity, thermal pressure, and ambipolar drift. Typically the sheet breaks up into a small number of fragments of elongated shape which are collapsing, losing magnetic flux through ambipolar drift, and interacting gravitationally with one another. The magnetic fields associated with these asymmetric clumps generate local vorticity (the net angular momentum of the sheet is identically zero). The characteristic vorticity structure is a vortex pair which flanks each clump and is associated with strong streaming motions along it. The clump masses are typically of order 1-10 $M_{\\odot}$, and scale with temperature like the Jeans mass, but are typically larger because of magnetic support. The main features of collapse in this geometry were predicted by the linear stability analysis of Zweibel (1998; hereafter Z98); there is also some overlap with the earlier stability analysis in 3D by \\citet{langer78}. Here we follow the evolution into the nonlinear regime and follow the growth of density fluctuations from .01 to up to 10 times the mean surface density. Collapse occurs on an intermediate time-scale, slower than the dynamical or free-fall time-scale, but faster than the ambipolar diffusion time-scale.  As noted above, this intermediate collapse rate may be observed in some clouds \\citep[{\\it e.g.} ][]{lm99,greg} . As no restrictions are placed on the morphology of the clumps, we can begin to explore the nature of flux loss and collapse in more complicated geometries than isolated axisymmetric clouds. One outcome suggested by linear theory which we have not resolved is whether stored magnetic energy is converted to turbulence, as there is no stored magnetic energy in the system initially.  The unperturbed initial geometry, governing equations, and linear theory are described in \\S\\ref{setup}.  Section \\ref{num} contains a description of the numerical method and main results: the collapse rate is discussed in \\S\\ref{gamgam}, the relationship between magnetic field {\\bf B} and density $\\rho$ or surface density $\\sigma$ in \\S\\ref{brho}, size of fragments in \\S\\ref{clump}, velocity structure in \\S\\ref{vel}, and distribution of energy in \\S\\ref{energy}. The validity of the approximations is discussed in \\S\\ref{discuss}, and the summary and conclusions are in \\S\\ref{summary}. ",
        "conclusions": "\\label{discuss}  It is difficult and computationally expensive to simulate the evolution of magnetized, self gravitating molecular clouds in three dimensions with sufficient resolution to capture all the relevant physical processes. In this study we focussed on the two-dimensional instabilities of a sheetlike cloud surrounded by a conducting medium without inertia. This allowed us to approximate the external magnetic field as current free, and to work in only two spatial dimensions. In the next two subsections we discuss the accuracy of these approximations.  \\subsection{Potential Field Approximation}\\label{potfield}  In the Appendix, we show how to extend each Fourier component of the magnetic and velocity perturbations above and below the sheet. Although the perturbations become highly nonlinear within the cloud, nonlinear effects outside the cloud are weak as long as the velocities inside the cloud are sub-Alfv\\'{e}nic with respect to the intercloud medium, which is expected for low ambient density. The compressive part of the cloud velocity field generates evanescent, fast magnetosonic waves which decay exponentially away from the disc, while the vortical part generates Alfv\\'{e}n waves which propagate away from the disc.  Both the magnetosonic and Alfv\\'{e}n waves slightly change the horizontal magnetic field perturbation in the disc, thereby changing the magnetic force from its value in the potential approximation. If we define an external Alfv\\'{e}n timescale $\\tau_{Ae}\\equiv (kv_{Ae})^{-1}$ for wavenumber $k$ in the disc, and let the timescale for the perturbation in the cloud be $\\tau_c$, then according to equation (\\ref{A9}) the correction to the force due to compressive motion is of order $(\\tau_{Ae}/\\tau_c)^2$, while the correction due to vortical motion is of order $(\\tau_{Ae}/\\tau_c){\\cal C}_0/{\\cal D}_0$, where ${\\cal C}_0/{\\cal D}_0$ is the ratio of the amplitude of compressive to noncompressive motion. If the motions in the disc were Alfv\\'{e}nic, $\\tau_{Ae}/\\tau_c$ would be of order the ratio of the cloud density to external density $(\\rho_c/\\rho_e)^{1/2}$, but the perturbation frequency is sub-Alfv\\'{e}nic, so the ratio of timescales is even larger. Moreover, the perturbations are primarily compressive rather than vortical. Thus, the error in the force incurred by the potential approximation is likely to be small. Even if the density contrast were only 10$^2$, and the motions in the disk were Alfv\\'{e}nic and purely vortical, the potential field approximation would still be accurate to 10\\%.  The Alfv\\'{e}n wave flux tends to suppress the instability, and removes vorticity from the cloud, at a rate that we can quantify. We define an energy damping time $\\gamma_d$ as the ratio of outward propagating energy flux to the vertically integrated wave energy in the disc. From equation (\\ref{A10}), \\begin{equation} \\label{eqn_damp} \\gamma_d = {{2\\rho_ev_{Ae}}\\over{\\sigma}} {{\\mid{\\cal C}_0^2\\mid}\\over{(\\mid{\\cal C}_0^2\\mid + \\mid{\\cal D}_0^2\\mid)}}. \\end{equation}  If we replace $\\sigma$ by the critical surface density $B_{0z}/2\\pi G^{1/2}$ and assume that the vertical fields inside and outside the cloud are the same then $\\gamma_d$ is just the gravitational frequency for the intercloud medium, reduced by the ratio of vortical to total kinetic energy \\begin{equation} \\gamma_d = (4\\pi G\\rho_e)^{1/2}{{\\mid{\\cal C}_0^2\\mid}\\over {(\\mid{\\cal C}_0^2\\mid + \\mid{\\cal D}_0^2\\mid)}}. \\end{equation}  Since self gravity is presumably negligible in the low density intercloud medium, the energy loss rate is negligible as well. Loss of vorticity is measured by the magnetic braking rate $\\gamma_{mb}$, which can be shown by a similar argument to be $\\gamma_{mb} = (4\\pi G\\rho_e)^{1/2}$.  We thus see how the potential field limit is approached as the density contrast between the cloud and intercloud medium increases. At the late stages of clump formation the potential field becomes highly distorted and develops partially closed topology, but we can ignore this for the relatively mild density contrasts studied in this paper.  \\subsection{Approximations to the Gas Physics}  The two dimensional approximation has a long and venerable history in galactic dynamics and accretion disc theory, as well as in studies of molecular clouds, and its errors for self gravitating systems are reasonably well understood. We expect the approximation to be reasonably good as long as the clump diameters exceed the disc thickness. The instability discussed in this paper has a 3D analog (Z98), but it must be treated by other means.  We assumed that the gas has an isothermal equation of state. This is a reasonable description of the kinetic pressure - density relationship, because of the high radiative efficiency of molecular gas. However, if the pressure were due to unresolved turbulence the medium would generally be less compressible; for example, Alfv\\'{e}n wave pressure follows density according to a 3/2 law \\citep{mz95}. This would make the medium more stable by increasing the value of $\\gamma_T$ (see eq. [\\ref{disp_rel}]), as would retention of magnetic pressure.  We took a uniform sheet at rest as an initial condition. This has the advantage of simplicity, but it means that there is no free energy stored in the background magnetic field. Thus, we have not tested the conjecture that the instability can convert magnetic free energy to turbulent energy, which was proposed in Z98.  We have treated ambipolar drift in the strong coupling approximation, and have implicitly assumed that $v_D <$ 20 km s$^{-1}$ (otherwise the rate coefficient would change). This is reasonable as long as the ion-neutral collision time, which is of order $5\\times 10^9 n_n^{-1}$ s, is less than other timescales in the problem.  We have parameterized the relationship between the collision rate and the surface density by an exponent $\\alpha$. In order to do better we would need a three dimensional model of the sheet and might need to follow the ionization as well. The results of linear theory are rather insensitive to the value of $\\alpha$, which suggests that it need not be calculated very accurately in the present models."
    },
    "0002/astro-ph0002217_arXiv.txt": {
        "abstract": "The time profiles of many gamma-ray bursts consist of distinct pulses, which offers the possibility of characterizing the temporal structure of these bursts using a relatively small set of pulse shape parameters.  This pulse decomposition analysis has previously been performed on a small sample of bright long bursts using binned data from BATSE, which comes in several data types, and on a sample of short bursts using the BATSE Time-Tagged Event (TTE) data type.  We have developed an interactive pulse-fitting program using the phenomenological pulse model of Norris, \\emph{et al.} and a maximum-likelihood fitting routine.  We have used this program to analyze the Time-to-Spill (TTS) data for all bursts observed by BATSE up through trigger number 2000, in all energy channels for which TTS data is available.  We present statistical information on the attributes of pulses comprising these bursts, including relations between pulse characteristics in different energy channels and the evolution of pulse characteristics through the course of a burst.  We carry out simulations to determine the biases that our procedures may introduce.  We find that pulses tend to have shorter rise times than decay times, and tend to be narrower and peak earlier at higher energies.  We also find that pulse brightness, pulse width, and pulse hardness ratios do not evolve monotonically within bursts, but that the ratios of pulse rise times to decay times tend to decrease with time within bursts. ",
        "introduction": "There has been considerable recent progress in the study of gamma-ray bursts.  Much of this results from the detection of bursts by BeppoSAX with good locations that have allowed the detection of counterparts at other wavelengths.  This has allowed measurements of redshifts that have firmly established that these bursts are at cosmological distances.  However, only a few redshifts are known, so there is still much work to be done in determining the mechanisms that produce gamma-ray bursts.  Investigation of time profiles and spectra can shed new light on this subject.  The vast majority of gamma-ray bursts that have been observed have been observed \\emph{only} by BATSE.  This data can be classified into three major types: burst locations with relatively large uncertainties, temporal characteristics, and spectral characteristics. Here, we shall examine temporal characteristics of bursts, along with some spectral characteristics.  The temporal structure of gamma-ray bursts exhibit very diverse morphologies, from single simple spikes to extremely complex structures.  So far, the only clear division of bursts based on temporal characteristics that has been found is the bimodal distribution of the $T_{90}$ and $T_{50}$ intervals, which are measures of burst durations~\\citep{kouveliotou:1993,meegan:1996}.  In order to characterize burst time profiles, it is useful to be able to describe them using a small number of parameters.  Many burst time profiles appear to be composed of a series of discrete, often overlapping, pulses, often with a \\emph{fast rise, exponential decay} (FRED) shape~\\citep{norris:1996}.  These pulses have durations ranging from a few milliseconds to several seconds. The different pulses might, for example, come from different spatial volumes in or near the burst source.  Therefore, it may be useful to decompose burst time profiles in terms of individual pulses, each of which rises from background to a maximum and then decays back to background levels.  Here, we have analyzed gamma-ray burst time profiles by representing them in terms of a finite number of pulses, each of which is described by a small number of parameters.  The BATSE data used for this purpose is described in Section~\\ref{sec:tts}.  The basic characteristics of the time profiles based on the above model are described in Section~\\ref{sec:results} and some of the correlations between these characteristics are described in Section~\\ref{sec:correlate}.  (Further analysis of these and other correlations and their significance are discussed in an accompanying paper, \\cite{lee:2000b}.)  Finally, a brief discussion is presented in Section~\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  Decomposing burst time profiles into a superposition of discrete pulses gives a compact representation that appears to contain their important features, so this seems to be a useful approach for analyzing their characteristics.  Our pulse decomposition analysis confirms a number of previously reported properties of gamma-ray burst time profiles using a larger sample of bursts with generally finer time resolution than in prior studies.  These properties include tendencies for the individual pulses comprising bursts to have shorter rise times than decay times; for the pulses to have shorter durations at higher energies; and for the pulses to peak earlier at higher energies, which is sometimes described as a hard-to-soft spectral evolution of individual pulses.  Pulse rise times tend to decrease during the course of a burst, while pulse decay times tend to increase.  When examining pulse widths, or durations, these two effects nearly balance each other; the apparent tendency for pulse widths to decrease during the course of a burst appears to be statistically insignificant.  The ratios of pulse rise times to decay times tend to decrease during the course of a burst. The evolution of pulse asymmetry ratios does not arise from selection effects in the pulse-fitting procedure, so it is most likely intrinsic to the bursters.  No other pulse characteristics show any time evolution within bursts, although it is possible that there is non-monotonic evolution; for example, a pulse characteristic may tend to be greater at the beginning and end of a burst and smaller in the middle, and the tests used here wouldn't be sensitive to this.  In particular, it doesn't appear that either pulse amplitudes or pulse count fluences have any tendency to increase or decrease during the course of a burst.  Also, later pulses in a burst don't tend to be spectrally harder or softer than earlier pulses, although there is spectral softening \\emph{within} most pulses.  The spectra of pulses within a burst also don't appear to be harder or softer for stronger or weaker pulses, or for longer or shorter duration pulses.  One may therefore conclude that the pulses in a burst arise from random and independent emission episodes such as those expected in the internal episodic shock model rather than the external shock models where the presence of distinguishable pulses must be attributed to inhomogeneities in the interaction of the blast wave shock and the clumpy interstellar medium.  When examining similar correlations between the attributes of some characteristic pulses from burst to burst, we find some weak and tantalizing evidence which may be due to cosmological redshift effects.  In the accompanying paper \\cite{lee:2000b} we describe the correlation studies which can distinguish between trends due to cosmological redshifts and intrinsic trends."
    },
    "0002/astro-ph0002021_arXiv.txt": {
        "abstract": "We present optical ($\\sim 3200$\\AA~ to $\\sim9000$\\AA~) off-nuclear spectra of 26 powerful active galaxies in the redshift range $0.1 \\leq z \\leq 0.3$, obtained with the Mayall and William Herschel 4-meter class telescopes. The sample consists of radio-quiet quasars, radio-loud quasars (all with $-23 \\geq M_{V} \\geq -26$) and radio galaxies of Fanaroff \\& Riley Type II (with extended radio luminosities and spectral indices comparable to those of the radio-loud quasars).  The spectra were all taken approximately 5 arcseconds off-nucleus, with offsets carefully selected so as to maximise the amount of galaxy light falling into the slit, whilst simultaneously minimising the amount of scattered nuclear light.  The majority of the resulting spectra appear to be dominated by the integrated stellar continuum of the underlying galaxies rather than by light from the non-stellar processes occurring in the active nuclei, and in many cases a 4000\\AA~break feature can be identified. The individual spectra are described in detail, and the importance of the various spectral components is discussed.  Stellar population synthesis modelling of the spectra will follow in a subsequent paper (Nolan et al. 2000). ",
        "introduction": "Understanding the host galaxies of active galactic nuclei (AGN) is now recognised as an important step on the path towards reaching an understanding of AGN themselves - how they form, how they are fuelled and how the differences between the various classes of object arise. In two areas in particular the nature of the host galaxy gives us a direct insight into the workings of the AGN: galaxy properties seem to play a role in determining the radio loudness of the central engine (powerful radio sources are almost never found in spiral or disc-dominated systems -- though see McHardy et al. 1994); and amongst radio-loud objects the host galaxies offer a powerful, orientation-independent means of testing models which attempt to unify different types of AGN via beaming and viewing angle effects ({\\it eg} Urry \\& Padovani 1995). The observations presented here attempt to address both of these issues by comparing the galaxies associated with the three main types of powerful AGN: radio-quiet quasars (RQQs), radio-loud quasars (RLQs) and radio galaxies (RGs) of Fanaroff \\& Riley Type II.  \\subsection{Quasar hosts}  Only once the problem of separating the diffuse galaxy emission from the wings of the quasar point spread function (PSF) has been overcome can one can begin to describe and classify the morphologies, brightness profiles and interaction histories of the quasar hosts. Over the last decade improvements in ground-based techniques and the advent of the Hubble Space Telescope (HST) have revolutionised our understanding of quasar host galaxies.  Evidence for mergers or interactions in the form of morphological disturbances and close companions is a common feature of these images, but a significant number of quasars are also found in what appear to be undisturbed hosts. In addition, the idea that radio-loudness is a straightforward consequence of the host galaxy type has had to be abandoned.  Although some radio-quiet quasars are found in spirals ({\\it eg} Hutchings et al. 1994; \\\"{O}rndahl, R\\\"{o}nnback \\& van Groningen 1997) in general the hosts of both radio-loud and radio-quiet quasars tend to have properties consistent with early-type galaxies ({\\it eg} V\\'{e}ron-Cetty \\& Woltjer 1990; Disney et al. 1995, Hutchings \\& Morris 1995, Bahcall et al. 1997, McLeod, Rieke \\& Storrie-Lombardi 1999) and typically have luminosities $> L^{*}$ ({\\it eg} Dunlop et al. 1993, Bahcall, Kirhakos \\& Schneider 1994, 1995ab, 1996; Hutchings et al. 1994, Boyce et al.  1998, Hooper, Impey \\& Foltz 1997) and in many cases are comparable in mass to brightest cluster galaxies (BCGs).  Meanwhile, McLeod \\& Rieke (1995a) find evidence for a {\\it lower} limit on the $H$-band luminosities (and hence the mass of the red, established stellar populations) of galaxies hosting radio-quiet AGN, which appears to increase as the nuclear luminosity increases, implying that the nuclear activity is closely linked to the mass of the bulge component of the host. This impression has been reinforced by McLure et al. (1999), who find that {\\it all} the RQQs with $M_{R}\\leq -24$ in their HST sample occur in massive elliptical galaxies, with only the least luminous radio-quiet objects lying in disc-dominated hosts.  Many long-held views about the triggering of nuclear activity and the origins of radio loudness are currently being reassessed in the light of these imaging studies, but images cannot tell the whole story. A completely independent way of characterising the host galaxies of AGN is via analysis and classification of their stellar populations and starformation histories. The aim of the observations described in this paper was to obtain high signal-to-noise spectra of quasar hosts and radio galaxies for use in spectrophotometric modelling to determine the nature and history of their stelllar constituents.  \\subsection{Spectroscopy of quasar hosts}  Previous off-nuclear spectroscopy of the host galaxies of quasars has produced mixed results. Boroson \\& Oke (1982) were the first to detect an unequivocably stellar continuum from the nebulosity surrounding the radio-loud quasar 3C48 and subsequent studies revealed stellar continua and emission/absorption features around several other quasars (Boroson, Oke \\& Green 1982; Boroson \\& Oke 1984; Boroson, Persson \\& Oke 1985; Hickson \\& Hutchings 1987; Hutchings \\& Crampton 1990). However, except for the general result that there appeared to be systematic differences between the spectra of radio-loud and radio-quiet quasar hosts, these studies produced little advance in our understanding of the relationship between RQQs, RLQs and RGs for the reasons outlined below.  In general the quasar targets were chosen virtually at random, often for the sole reason that they were `interesting' and/or unusual. Until now a programme of off-nuclear spectroscopy for statistically useful and properly matched samples of RQQs and RLQs has never been carried out, nor has any attempt been made to compare the off-nuclear spectra of RLQs with equivalent {\\it off-nuclear} spectra of radio galaxies.  The early work also focussed mostly on emission-line activity, with discussion of the stellar continuum being confined to classification as red or blue, and the identification of a few stellar features. Only limited attempts were made to use the form of the spectrum to investigate the composition and evolution of the stellar population.  Finally, few of the host galaxy spectra were taken sufficiently off-nucleus - {\\it eg} the spectra taken by Boroson and collaborators were taken only 3$''$ from the quasar. This was done in the belief that any further off-nucleus the host galaxy would be too faint for a reasonable spectrum to be obtained, but inevitably resulted in significant contamination of the host galaxy spectrum by scattered light from the quasar nucleus. A scaled version of the quasar spectrum had therefore to be subtracted from the off-nuclear spectrum to reveal the spectrum of the underlying host, but the extra noise and systematic errors introduced by this process severely limited the quality of the final spectra.  Thus, the main obstacles in these previous attempts to classify the hosts of powerful AGN were the difficulty in separating the underlying starlight from the glare of the quasar and - less immediately, but still of some importance - the lack of well-defined samples of a sufficient size to carry out straightforward statistical comparisons.  As far as possible, the design of the current study was intended to circumvent both these problems with the goal of obtaining clean spectra of the stellar component of radio-loud and radio-quiet quasar hosts and radio galaxies, which could then be used to search for systematic differences and similarities between the stellar populations and of the galaxies hosting each type of activity. Despite that fact that the radio galaxies lack bright nuclear point sources we were careful to adopt exactly the same observing strategy as used for the quasars to ensure that the data would be directly comparable.  In this paper we describe the observations and present the spectra.  A second paper (Nolan et al. 2000) will describe the results of spectrophotometric modelling to estimate the ages and starformation histories of the galaxies (preliminary results have already been reported by Kukula et al. 1997).  The current paper is organised as follows. Section~2 describes the samples used and Section~3 details the observing and data reduction strategies chosen to optimise the amount of starlight collected by the various instruments employed throughout the study. Section~4 gives an overview of the data obtained and Sections~5 and 6 contain more detailed information on the individual spectra. ",
        "conclusions": "We describe and present optical spectra of 26 galaxies hosting powerful nuclear activity. The sample contains carefully matched subsamples of all three types of powerful AGN - radio-quiet and radio-loud quasars, and FR{\\sc ii} radio galaxies - enabling us to investigate the relationship between the host galaxy and the radio properties of the resident AGN and also to test unified models for radio-loud objects.  The spectra were taken 5~arcsec {\\it off-nucleus} and, via a careful choice of slit position, aim to maximise the amount of galaxy light entering the instrument whilst avoiding contamination from the active nucleus. In the majority of cases this approach appears to have been successful; the continuum is clearly overwhelmingly stellar in origin and even when the presence of scattered nuclear emission is suspected, features such as the 4000\\AA~break are still clearly discernable.  In almost all objects we detect at least some stellar signatures.  A second paper (Nolan et al. 2000) will discuss spectrophotometric modelling of the galaxy spectra in order to determine the ages and starformation histories of their constituent stellar populations."
    },
    "0002/astro-ph0002398_arXiv.txt": {
        "abstract": "\\small  We have observed the OH (1720 MHz) line in five galactic SNRs with the VLA to measure their magnetic field strengths using the Zeeman effect.  We detected all 12 of the bright ($S_{\\nu} > 200$ mJy) OH (1720 MHz) masers previously detected by \\markcite{Fra96}Frail et al.  (1996) and \\markcite{Gre}Green et al.  (1997) and measured significant magnetic fields (i.e.  $ > 3\\sigma$) in ten of them.  Assuming that the ``thermal'' Zeeman equation can be used to estimate $\\mid\\vec{B}\\mid$ for OH masers, our estimated fields range from 0.2 to 2 mG.  These magnetic field strengths are consistent with the hypothesis that ambient molecular cloud magnetic fields are compressed via the SNR shock to the observed values.  Magnetic fields of this magnitude exert a considerable influence on the properties of the cloud with the magnetic pressures ($10^{-7} - 10^{-9}$ erg \\cc\\/) exceeding the pressure in the ISM or even the thermal pressure of the hot gas interior to the remnant.  This study brings the number of galactic SNRs with OH (1720 MHz) Zeeman detections to ten. ",
        "introduction": "Magnetic fields can moderate the impact that a shock has on a molecular cloud.  In the absence of a field, a supernova blast wave will heat, compress and fragment the cloud and may ultimately destroy it (\\markcite{Kle}Klein, McKee \\& Colella 1994).  The inclusion of a field mitigates these effects, limiting compression and stabilizing it against fragmentation, allowing the cloud to survive, perhaps to trigger a future generation of star formation (\\markcite{Mie}Miesch \\& Zweibel 1994; \\markcite{MacL}MacLow et al.  1994).  It has long been known that supernova remnants (SNRs) possess magnetic fields.  Observations of synchrotron radiation have established that the {\\it direction} of the magnetic field in young SNRs like Cas A is predominately parallel to the shock normal (i.e.  radial), whereas for older remnants the fields are perpendicular (e.g.  \\markcite{Dic}Dickel et al.  1991; \\markcite{Mil}Milne 1990).  The latter geometry likely originates from the compression of the ambient interstellar field, while Rayleigh-Taylor instabilities are invoked to explain the radial fields in young SNRs (\\markcite{Jun}Jun \\& Norman 1996).  Until recently, estimates of the {\\it strength} of the magnetic fields in SNRs had to rely on the somewhat dubious equipartition approximation.  This situation has changed with the re-discovery of shock-excited OH (1720 MHz) maser emission in SNRs (\\markcite{Fra94}Frail, Goss \\& Slysh 1994).  In a series of recent papers, the satellite line of the OH molecule at 1720.53 MHz has been used as a powerful probe of SNR-molecular cloud interactions.  OH (1720 MHz) masers are found in $\\sim $20 SNRs, or 10\\% of the known SNRs in our Galaxy (see \\markcite{Kor}Koralesky et al. 1998 and references therein).  This maser line is inverted through collisions with H$_2$ ($n\\sim$~few$\\times{10}^4$ cm$^{-3}$ and T$\\sim $80 K) behind C-type shocks propagating into molecular clouds (\\markcite{Rea99}Reach \\& Rho 1999, \\markcite{Fra98}Frail \\& Mitchell 1998). The geometry of the shock is well constrained since strong maser amplification can only occur when the shock front is viewed edge on (\\markcite{Cla97}Claussen et al.  1997).  These observational statements are well-supported by theoretical modeling of the OH (1720 MHz) excitation (\\markcite{Loc}Lockett, Gauthier \\& Elitzur 1998; \\markcite{War99}Wardle, Yusef-Zadeh, \\& Geballe 1999; \\markcite{Wardle}Wardle 1999).  One advantage of observing the OH (1720 MHz) maser line is that it allows for a measurement of the strength of the magnetic field via Zeeman splitting (e.g.  \\markcite{Tro}Troland \\& Heiles 1982).  In this case, when the observed splitting is small compared to the line width, V $= ZC$\\Bvm\\/\\thinspace{dI/d$\\nu$}, where \\Bvm\\/ is the the total magnetic field strength, Z is the Zeeman splitting coefficient, and $C$ is a constant which depends on the angle between the line of sight and $\\vec{B}$ (the possible forms of $C$ will be discussed in \\S 4.1).  For now we will denote the combination of $C$\\Bvm\\/ as \\Bth\\/.  This method has already been used to successfully measure the magnetic field strength in the post-shock gas behind five SNRs (Sgr A East, W44, W28, G32.8$-0.1$, and G346.6$-$0.2), yielding values for \\Bth\\/ between 0.1 to 4 mG (\\markcite{Yus96}Yusef-Zadeh et al.  1996, \\markcite{Cla97}Claussen et al.  1997, \\markcite{Kor}Koralesky et al.  1998).  With these promising results in mind, we have performed 1720 MHz Zeeman studies toward five more of the 17 SNRs found in surveys by \\markcite{Fra96}Frail et al.  (1996) and \\markcite{Gre}Green et al.  (1997) to contain OH (1720 MHz) masers using the NRAO\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under a cooperative agreement by Associated Universities, Inc.}  VLA.  Only SNRs with bright masers (I$>$200 mJy) and no previous high resolution Stokes V observations were chosen. ",
        "conclusions": "We have observed the OH (1720 MHz) line in five galactic SNRs to measure their magnetic field strengths using the Zeeman effect.  We detected all 12 of the bright ($S_{\\nu} > 200$ mJy) OH (1720 MHz) masers previously observed by \\markcite{Fra96}Frail et al.  (1996) and \\markcite{Gre}Green et al.  (1997) and measured significant magnetic fields (i.e.  $ > 3\\sigma$) in ten of them.  The estimated field strengths, which we denote \\Bth\\/, range from 0.2 to 2 mG and are in good agreement with those measured in the five other SNR (0.1 - 4 mG) for which Zeeman OH (1720 MHz) maser studies exist (Koralesky et al. 1997 and references therein).  In these studies, the field strengths were calculated using the thermal Zeeman equation for Zeeman splitting less $<<$ the line width (i.e.  V $=Z$\\Bth\\/\\thinspace{$d$I$/d\\nu$}).  In this formula, the total field strength (\\Bvm\\/) is related to \\Bth\\/ by \\Bvm\\/=\\Bth\\//$C_m$ where $C_m$ is a function of the viewing angle $\\theta$ ($C_m=cos\\thinspace{\\theta}$ for thermal emission).  However, there exists some controversy about the $\\theta$ dependence of $C_m$ for saturated masers and whether it is even a constant for unsaturated masers (see \\S 4.1).  For example, in the Elitzur (1998) model $C_m\\propto 1/(cos\\thinspace{\\theta})$ for masers.  {\\em From comparison of these data with maser polarization studies we conclude that these OH (1720 MHz) masers are likely to be saturated and that $C_m\\lesssim 5$, i.e.  the \\Bth\\/ values reported in Table 3, overestimate the actual field strengths (\\Bvm\\/) by less than a factor of five.}  We estimate that these magnetic field strengths are consistent with the hypothesis that ambient molecular cloud B-fields are compressed via the SNR shock to the observed values (\\S 4.2).  Indeed, field strengths of this magnitude exert a considerable influence on the properties of the shocked cloud.  For example the magnetic pressures estimated from the values of \\Bth\\/ listed in Table 3 ($10^{-7} - 10^{-9}$ erg \\cc\\/) far exceed the pressure in the ISM or even the thermal pressure of the hot gas interior to the remnant (\\S 4.3).  We also find that this range of magnetic pressures is in very good agreement with the ram pressure expected from C-type radiative shocks.  In \\S 4.2 and \\S 4.3 we show that there is excellent agreement between our values of \\Bth\\/ and those implied by shock compression and ram pressure using typical values from the literature. It is somewhat difficult to maintain this agreement if \\Bth\\/ should be reduced by as much as a factor of $\\lesssim 5$ (i.e.  Elitzur's maser polarization model \\S 4.1), although there is probably sufficient uncertainty in the parameters to allow such fine-tuning.  We also show that the observed values of \\Bth\\/ for OH (1720 MHz) masers are greater than those observed in molecular clouds for the same range of densities (\\markcite{Cru}Crutcher 1999).  It is possible that this difference is due the intrinsic physical nature of the $\\mid\\vec{B}\\mid \\propto n^{1/2}$ relation (\\S 4.2) or may be the result of overestimating the maser B-fields (\\S 4.1).  In the future, knowledge of the field strength could be used in conjunction with molecular data to study the physics of these molecular shocks in more detail.  In one such study \\markcite{Fra98}Frail \\& Mitchell (1998) mapped the distribution of molecular gas in the vicinity of several masers spots in W28 and W44 with the JCMT.  These observations revealed that the OH masers are preferentially located along the edges of thin filaments or clumps of molecular gas, suggesting compression of the gas by the passing shock.  In addition to this morphological evidence, the density, temperature and velocity dispersion of the gas at these locations suggested that the OH (1720 MHz) masers originate in post-shock gas.  Combining the VLA and JCMT data they were able to show directly that the magnetic pressure dominates over the thermal pressure in the post-shock gas, balancing against the ram pressure of the gas entering the shock wave (i.e.  B$_{ps}^2/8\\pi=\\rho_\\circ{\\rm V}_s^2$).  Thus OH (1720 MHz) measurements of \\Bth\\/ lead to constraints on the physics of these molecular shocks that are difficult to obtain any other way.  When used in conjunction with molecular observations, it should be possible to fully specify the properties (i.e.  geometry, density, temperature, velocity) of C-type shocks."
    },
    "0002/astro-ph0002351_arXiv.txt": {
        "abstract": "The observations of light elements in the Sun and Solar type stars give special clues for understanding the hydrodynamical processes at work in stellar interiors. In the Sun $^7$Li is depleted by 140 while $^3$He has not increased by more than  $\\cong~10\\%$ in 3 Gyrs. Meanwhile the inversion of helioseismic modes lead to a precision on the sound velocity of about $.1\\%$. The mixing processes below the solar convection zone are constrained by these observations. Lithium is depleted in most Pop~I solar type stars. In halo stars however, the lithium abundance seems constant in the ``spite plateau\" with no observed dispersion, which is difficult to reconcile with the theory of diffusion processes. In the present paper, the various relevant observations will be discussed. It will be shown that the $\\displaystyle \\mu $-gradients induced by element settling may help solving the ``lithium paradox\". ",
        "introduction": "Element diffusion and mixing processes in stellar interiors are now widely constrained, first by detailed observations of abundances, second by helio and asteroseismology. In most cases however, pure microscopic diffusion in stars would lead to abundance variations much larger than those observed : mild macroscopic motions in stellar radiative zones are definitely needed to account for the observations. This gives strong constraints on the kind of mixing processes allowed. Other constraints come from the consequences of the nuclear reactions occuring in stellar interiors : in some cases stellar mixing from the atmosphere down to the regions of nuclear processing is needed to explain the observed element abundances. This is the case, for example, to account for the depletion of lithium in the Sun and solar type stars.  Lithium observations in main-sequence population I field stars and galactic clusters show a large abundance dispersion which has been extensively studied in the literature (see reviews by Deliyannis 2000, Charbonnel 2000, Michaud 2000 and Pinsonneault 2000). The lithium abundance decreases for decreasing effective temperature below 5500K and the depletion increases with increasing age. This is generally attributed to the deepening of the convective zone, associated with some mild mixing process connecting the bottom of the convective zone with the nuclear destruction region.  Lithium is also depleted in F-type stars (the so-called ``Boesgaard dip\"). Several possible reasons have been invoked to explain this feature, most related to mixing and nuclear destruction. Element segregation has been proved negligible here as it would lead to unobserved variations of metal abundances (Turcotte et al 1998) and beryllium (Boesgaard 2000).  On the other hand, observations of lithium in main-sequence population~II field stars show remarkably constant abundances, with a very small dispersion (e.g. Bonifacio and Molaro 1997) Why is lithium destroyed in Pop~I stars while it does not seem destroyed in Pop~II stars?  For the same effective temperatures, the convective zone is smaller in Pop~II stars than in Pop~I stars because of their smaller metallicity. Meanwhile they have a smaller rotation velocity on the average. This could explain why the lithium destruction induced by nuclear reactions is smaller in these stars than in Pop~I stars. However the element segregation is more important for smaller densities and smaller rotation, so that this process should lead to a visible lithium depletion, which is not observed (Vauclair and Charbonnel 1995 and 1998). This represents the so-called ``lithium paradox\". Here we suggest that the influence of $\\displaystyle \\mu $-gradients on the rotation-induced mixing may help solving this paradox. ",
        "conclusions": "There are many observations in stars which give evidences of mixing processes occuring below the outer convective zones as, for example, the lithium depletion observed in the Sun and in galactic clusters. The process we have described above should not apply in all these stars. The reason could be related to the rapid rotation of young stars on the ZAMS and to their subsequent rotational braking.   The abundance determinations in the solar photosphere show that lithium has been depleted by a factor of about 140 compared to the protosolar value while beryllium has not been depleted by more than a factor 2, and maybe much less, as discussed by Balachandran and Bell (1997). These values represent strong constraints on the mixing processes in the solar interior.  Observations of the $^3$He/$^4$He ratio in the solar wind and in the lunar rocks (Geiss 1993, Geiss and Gloecker 1998) show that this ratio may not have increased by more than $\\cong~10\\%$ since 3 Gyr in the Sun. While the occurence of some mild mixing below the solar convective zone is needed to explain the lithium depletion , the $^3$He/$^4$He observations put a strict constraint on its efficiency. The only way to obtain such a result is to postulate a mild mixing, which would be efficient down to the lithium nuclear burning region but not too far below, to preserve the original $^3$He abundance. The efficiency of this mixing should also decrease with time, as the $^3$He peak itself builts up during the solar life.  It is interesting to compute the minimum enhancement of the $^3$He/$^4$He ratio implied by the lithium observed depletion. Vauclair and Richard 1998 showed that it is possible to deplete lithium by a factor larger than $100$ as observed and not increase $^3$He/$^4$He by more than 5 percent since the solar origin.  In this case beryllium is only depleted by about 10 percent.  Such a confined mixing zone is also needed from helioseismology~: although the introduction of pure element settling in the solar models considerably improves the consistency with the seismic Sun, some discrepancies do remain, particularly below the convective zone where a \"spike\" appears in the sound velocity (Richard et al 1996, Turck-Chi\\`eze et al. 1998). It has been shown that this behavior may be due to the helium gradient which would be too strong in case of pure settling. Mild macroscopic motions below the convective zone slightly decrease this gradient and helps reducing the discrepancy (Richard et al 1996, Corbard et al 1998, Brun et al 1998). The helium profiles directly obtained from helioseismology (Basu 1998, Antia and Chitre 1998) show indeed a helium gradient smoother than the gradient obtained with pure settling.  The constraints implied by both the helioseismic inversions and abundance determinations in the Sun converge towards the existence of a small mild mixing region below the convective zone, which would extend down to a depth of the order of one scale height. The implied mixing region must be very mild, with diffusion coefficients of $10^3$ - $10^4$ only. It must also be completely deconnected from the solar core. No mixing can indeed be allowed down to the nuclear energy production region as it would lead to a sound velocity incompatible with helioseismology. In particular the mixing processes invoked by Morel and Schatzman 1996 to decrease the neutrino fluxes are excluded by helioseismology (Richard and Vauclair 1997).  Mixing processes localized at the boundary between convective and radiative regions include overshooting and regions of large differential rotation like the ``tachocline\" below the solar convective zone. Up to now, overshooting was generally treated in the models simply as a continuation of the convective zone on a fraction of a pressure scale height. Recent parametrisations use a diffusion coefficient which decreases exponentially with decreasing radius (Freytag et al 1996). The tachocline, which represents in the present Sun the small boundary between the region of large differential rotation (in the convective zone) and the region of solid rotation (in the radiative zone below) is also treated as a mixed layer with an exponentially decreasing diffusion coefficient (Brun et al 1998, Richard 1999). Results are encouraging, although more sophisticated numerical simulation including 2-D abundance variations would be needed to go further.  In any case, the self-regulating process that we have discussed for halo stars in section~3 would not apply below the convective zone in the Sun and solar type stars because of the differential rotation which takes place there. Such a differential rotation would not be expected in halo stars if we suppose that they always rotated slowly and thus did not suffer large transport of angular momentum. The different behavior for the lithium abundance in Pop~I and Pop~II stars could thus be directly related to their rotation history."
    },
    "0002/astro-ph0002167_arXiv.txt": {
        "abstract": "The Advanced Satellite for Cosmology and Astrophysics (ASCA) has made multiple observations of the Small Magellanic Cloud (SMC). X-ray mosaic images in the soft (0.7--2.0~keV) and hard (2.0--7.0~keV) bands are separately constructed, and the latter provides the first hard X-ray view of the SMC. We extract 39 sources from the two-band images with a criterion of $S/N>5$, and conduct timing and spectral analyses for all of these sources. Coherent pulsations are detected from 12 X-ray sources; five of which are new discoveries. Most of the 12 X-ray pulsars are found to exhibit long-term flux variabilities, hence they are likely to be X-ray binary pulsars (XBPs). On the other hand, we classify four supernova remnants (SNRs) as thermal SNRs, because their spectra exhibit emission lines from highly ionized atoms. We find that XBPs and thermal SNRs in the SMC can be clearly separated by their hardness ratio (the ratio of the count rate between the hard and soft bands). Using this empirical grouping, we find many XBP candidates in the SMC, although no pulsations have yet been detected from these sources. Possible implications on the star-formation history and evolution of the SMC are presented by a comparison of the source populations in the SMC and our Galaxy. ",
        "introduction": "Bright X-ray sources such as supernova remnants (SNRs) and binaries with a neutron star (NS) or a black hole component are `relics' of massive stars and hence carry key information on the star formation history, dynamics, and structure of their host galaxy. However, our knowledge of the populations of bright X-ray sources in our Galaxy is highly incomplete due to the Galactic absorption in the low energy band, large angular size that must be observed, and limited information on the distance of the relevant objects.   The Small Magellanic Cloud (SMC), a satellite of our Galaxy, is the next nearest neighbor after the Large Magellanic Cloud (LMC). The proximity (60~kpc is assumed in this paper; Mathewson 1985\\markcite{Mathewson1985a}), reasonable angular size ($\\sim 3^\\circ \\times 3^\\circ$), and low interstellar absorption to the SMC are all favorable for an unbiased survey of the X-ray source populations in the entire galaxy. Surveys of soft X-ray sources (below $\\sim 2$~keV) have been carried out with the Einstein and ROSAT satellites. The most complete source catalogs are presented in Wang \\& Wu (1992, hereafter WW92\\markcite{Wang1992}) and Kahabka et al. (1999, hereafter K99\\markcite{Kahabka1999}). K99 classified the sources detected by the ROSAT PSPC (Position Sensitive Proportional Counter), based on two different hardness ratios, together with spatial information. Cowley et al. (1997)\\markcite{Cowley1997} and Schmidtke et al. (1999)\\markcite{Schmidtke1999} made numerous observations of the SMC with the ROSAT HRI (High Resolution Imager), and determined accurate positions of X-ray sources. Optical observations made at CTIO (Cerro Tololo Interamerican Observatory) showed that five of these X-ray sources have a massive star companion and hence are high mass X-ray binaries (HMXBs).   Bright X-ray sources can be classified by their spectral and temporal characteristics. X-rays from young SNRs are either due to a shock heated plasma (in shell-like SNRs) which preferentially emits soft (below $\\sim 2$~keV) line-dominated X-rays, or due to the release of rotational energy from isolated NSs (X-ray pulsars associated with Crab-like SNRs) with hard and featureless spectra. Unlike SNRs, the other classes of X-ray sources emit time-variable X-rays. NSs with a low mass companion star (low mass X-ray binaries: LMXBs) have hard spectra and often exhibit X-ray bursts. Black hole binaries (BHs) have bimodal behavior and show both high-soft and low-hard spectral states. NSs with a high mass stellar companion (HMXBs) have even harder X-ray spectra and often exhibit long-term flux variations of a factor $\\gtrsim 10$--100. Many HMXBs also show coherent pulsations (X-ray binary pulsars; XBPs). Consequently, most bright X-ray sources emit copious hard X-rays (above $\\sim 2$ keV). However, due mainly to a lack of imaging capability in the hard X-ray band, no systematic survey of the populations of hard X-ray sources had previously been made. The Advanced Satellite for Cosmology and Astrophysics (ASCA; Tanaka, Inoue, \\& Holt 1993\\markcite{Tanaka1993}) is equipped with wide-band X-ray imaging instruments and hence is optimized for an X-ray population study.  In this paper, we report the results from multiple observations on the SMC with ASCA. The observation fields and the method of data reduction are presented in \\S\\ref{sec:obs}. In \\S\\ref{sec:result}, we extract X-ray sources detected above the 5-$\\sigma$ level from the observation fields, perform temporal and spectral analyses on all the sources, and categorize them into thermal SNRs and XBPs. \\S\\ref{sec:comments} gives details of individual X-ray pulsars and SNRs. In \\S\\ref{sec:dis}, we propose a simple and reliable source classification method. We clearly distinguish between thermal SNRs and XBPs with this method, then discuss these X-ray source populations in the SMC. Finally, we summarize this study in \\S\\ref{sec:conc}. ",
        "conclusions": "}  We have performed systematic analyses of X-ray sources in the SMC with all the available data obtained with ASCA.  Among the 16 X-ray pulsars in the SMC, five were newly discovered during our analyses. We also confirmed the pulsations of another seven pulsars. We regarded 14 pulsars out of the 16 as XBPs because of their flux variability and, for some sources, the existence of an optical counterpart.  Eight radio SNRs were detected. The spectral analyses revealed that at least four SNRs exhibit the emission lines of ionized atoms, hence they are expected to be thermal SNRs. Another two SNRs are possibly emitting non-thermal X-rays.   We established a simple and reliable method of source classification using the hardness ratio (HR). XBPs fall in a range of $0.2\\leq {\\rm HR} \\leq0.6$, while thermal SNRs fall in a range of $-1\\leq {\\rm HR}\\leq -0.6$. With this method, eight XBP candidates and one thermal SNR candidate are found. Normalized by the galaxy mass, HMXBs are found to be extremely populous in the SMC, compared with those in our Galaxy. An even more striking contrast between the two galaxies is the number ratio of HMXBs to LMXB; it is 33/0 in the SMC but is 70/100 in our Galaxy. The lack of Crab-like SNRs in the SMC may indicate that the formation of single massive stars in the SMC was less active than that of binary systems."
    },
    "0002/astro-ph0002484_arXiv.txt": {
        "abstract": "We have mapped a giant molecular cloud in the giant HII region NGC\\,604 in M33 in the 492 GHz $^3P_1 \\rightarrow ^3P_0$ transition of neutral atomic carbon using the James Clerk Maxwell Telescope.  We find the distribution of the [CI] emission to be asymmetric with respect to the CO J=$1\\to0$ emission, with the peak of the [CI] emission offset towards the direction of the center of the HII region.  In addition, the line ratio I$_{[CI]}$/I$_{CO}$ is highest ($\\sim$ 0.2) facing the HII region and lowest ($\\lesssim$ 0.1) away from it. These asymmetries indicate an edge-on morphology where the [CI] emission is strongest on the side of the cloud facing the center of the HII region, and not detected at all on the opposite side  This suggests that the sources of the incident flux creating C from the dissociation of CO are the massive stars of the HII region.  The lowest line ratios are similar to what is observed in Galactic molecular clouds,  while the highest are similar to starburst galaxies and other regions of intense star formation.  The column density ratio, N(C)/N(H$_2$) is a few $\\times$~10$^{-6}$, in general agreement with models of photodissociation regions. ",
        "introduction": "The fine structure line of atomic carbon at 492 GHz (607 $\\mu$m) is one of the important cooling lines of photon dominated regions (PDRs). PDRs form the transition zones between the molecular interstellar medium (ISM) and the atomic ISM, where molecular gas is dissociated by ultraviolet photons from nearby star forming regions.  Maps of the distribution of [CI] emission in molecular clouds and on larger scales in nearby galaxies have revealed that it is often cospatial with the emission from CO (e.g., Keene et al. 1985; Israel, White \\& Baas 1995), despite the fact that the ionization energy of carbon is very close to the photodissociation energy of CO. From these energy considerations, [CI] would be expected to exist only in a thin layer at the surface of molecular clouds.  The similarity between the spatial distributions of CO and [CI] is evidence for a clumpy molecular ISM that allows ultraviolet photons to penetrate into the interiors of molecular clouds (White \\& Padman 1991).  The ratio of atomic carbon to CO, N(C)/N(CO), depends on which chemical reactions occur in the PDR, which in turn depends upon the degree of ionization of the gas.  When the ionization fraction exceeds a critical value, the gas phase chemistry is dominated by charge transfer reactions with H$^+$. These reactions suppress the amount of H$_3^+$, the most important species for reactions that destroy atomic carbon. Thus, at high ionization fractions, the abundance of C can remain high, while at lower ionization fractions H$_3^+$ reactions dominate and C is destroyed (Graedel, Langer \\& Frerking et al. 1982; Flower et al. 1994).  Both UV photons and cosmic rays can influence the ionization fraction.  A strong UV field incident upon a molecular cloud will eventually be stopped by the increasing optical depth as the radiation penetrates the cloud. Thus, the distribution of [CI] emission should appear asymmetric.  The flux of energetic cosmic rays, on the other hand, will be isotropic within a galaxy and the cosmic rays will be capable of penetrating the entire molecular cloud.  If the cosmic rays are a more important source of ionization than the UV flux, the [CI] distribution should appear similar to the CO distribution, with both species tracing the overall distribution of carbon in the ISM.  We present observations of the [CI] line at 492 GHz of the giant molecular cloud (GMC) NGC\\,604-2, in the giant HII region NGC\\,604, which is in the Local Group galaxy M33.  By observing a GMC in a nearby galaxy (D = 0.84 Mpc; Freedman, Wilson, \\& Madore 1991), we can map the extent of the cloud in a comparatively short time, and thus determine the spatial distribution of the [CI] emission.  NGC\\,604-2 has a molecular mass of $6.3~\\times~10^5$ M$\\sun$, a diameter of $\\sim$ 32 pc, and a linewidth of $\\sim$ 11 \\kms\\thinspace (Wilson \\& Scoville 1992).  It is similar to the larger GMCs in the Milky Way Galaxy. NGC\\,604 is the most luminous HII region in M33, and is the second nearest giant HII region, the nearest being 30 Doradus in the Large Magellanic Cloud.  Although these two HII regions are comparable in size, the H$\\alpha$ luminosity of NGC\\,604 is approximately half that of 30 Doradus (Kennicutt 1984).  The most important difference between the two regions is the distribution of their ionizing stars: 30 Doradus is dominated by the R136 cluster (Walborn 1991), while in NGC\\,604 the distribution of massive stars is not so concentrated towards the center of the HII region (Drissen et al. 1993).  Hunter et al. (1996) have calculated the average density of O stars in NGC\\, 604 to be 0.0018 stars pc$^{-2}$, 100 times lower than in R136 and more than a factor of 2 lower than the average density over the entire 30 Doradus region. ",
        "conclusions": "\\subsection{Comparison to Galactic Molecular Clouds}  The best comparison to our observations of NGC\\,604-2 is the sample of four galactic molecular clouds observed by Plume et al. (1999).  They mapped the clouds in several lines, including [CI] and $^{13}$CO J=$2\\to1$. Like us, they mapped most of each cloud, though they obtained a higher physical resolution since their clouds are galactic.  \\subsubsection{Line Ratios}  Plume et al. obtained line ratios of I$_{[CI]}$/I$_{^{13}CO J=2\\to1}$ at various positions in their clouds, with an average line ratio of 0.59 $\\pm$ 0.19.  For comparison with our data we convert this to I$_{[CI]}$/I$_{CO}$ (where I$_{CO}$ refers to the $^{12}$CO J=$1\\to0$ line) using the average line ratio for $^{12}$CO J=$2\\to1$ to $^{12}$CO J=$1\\to0$ of 0.7 for the Orion A and B clouds from Sakamoto et al. (1994), and the average value of the $^{12}$CO to $^{13}$CO J=$2\\to1$ line ratio of 5.5 $\\pm$ 1 measured for the Milky Way by Sanders et al.(1993). The Plume et al. value of I$_{[CI]} $/I$_{CO}$ becomes 0.08 $\\pm$ 0.03, slightly lower than would be allowed by the scatter about our mean value, but consistent within the error bars with some of our individual positions.  If we adopt different, but still reasonable, values for the two line ratios, the results agree with our observations better still.  Using the $^{12}$CO J=$2\\to1$ to $^{12}$CO J=$1\\to0$ line ratio of 1.3 found near HII regions in Orion instead of the global average, and using the $^{12}$CO to $^{13}$CO J=$2\\to1$ ratio of 4.5 $\\pm$ 0.7 found by Wilson, Howe \\& Balogh (1999) for the molecular cloud M17, we find I$_{[CI]}$/I$_{CO}$ = 0.17 $\\pm$ 0.06, in good agreement with our observations.  Clearly because Plume et al. did not observe $^{12}$CO J=$1\\to0$, comparisons between our data and theirs will be inexact.  We do, however, find a general consistency.  On the edges of the molecular clouds, where the CO is less shielded, Plume et al. find values of the line ratio up to 7 times higher.  We see a similar trend in NGC\\,604-2, but our physical resolution at the distance of M33 is 50 pc, compared to the 1 - 2 pc for Plume et al.  Thus, we are averaging over large areas within the cloud and smoothing away any large contrasts in the line ratio that may exist. Observations with a submillimeter interferometer would help in this case.  Plume et al. find a strong correlation between [CI] line strength and $^{13}$CO J=$2\\to1$ line strength, and fit both a power-law and linear relation to their data.  We can use the $^{13}$CO J=$2\\to1$ observation by Wilson, Walker \\& Thornley (1997) to determine if NGC~604-2 follows either relation.  The beam size of Wilson et al. is 22$\\arcsec$, so we scale their $^{13}$CO J=$2\\to1$ intensity by (22)$^2 / $(12)$^2$ to make an approximate correction for the different beam sizes.  Using this value in the two fitted relations, we find that either of them reproduces our I$_{[CI]}$ observation at the (0,0) position. The linear relation gives I$_{[CI]}$ = 2.7 $\\pm$ 0.2, compared to the measured value of 2.8 $\\pm$ 0.3  The power law gives 3.0 $\\pm$ 0.2. Our data point falls close to both relations in the plot of I$_{[CI]}$/I$_{^{13}CO J=2\\to1}$ versus I$_{^{13}CO J=2\\to1}$ in Plume et al., their Figure~10.  Both our data point, and the data of Plume et al. are in qualitative agreement with single layer PDR models.  \\subsubsection{Column Densities}  We obtain column densities a factor of 3 to 7 smaller than the typical column densities measured by Plume et al. The difference may be explained by the lower than solar metallicity of NGC\\,604, or the fact that NGC\\,604-2 lies in a giant HII region, compared to the more normal galactic star forming regions of the Plume et al. clouds.  They find lower column densities at the edges of their clouds, and higher densities at the cloud cores, which we do not, but this is again likely due to our lower physical resolution.  We compare the column density ratio N(C)/N(H$_2$) in NGC\\,604-2 with the observations by Plume et al. (1999) and predictions of models by Flower et al. (1994) and Spaans \\& van Dishoeck (1997).  Plume et al. calculate densities of C and $^{13}$CO, and the theoretical papers derive column densities of C, $^{12}$CO, and $^{13}$CO.  To convert these CO densities to H$_2$ densities, we use a $^{13}$CO/H$_2$ ratio of 1.5~$\\times~10^{-6}$ (Bachiller \\& Cernicharo 1986) and an isotope abundance ratio $^{12}$C/$^{13}$C of 62 (Langer \\& Penzias 1993).  These numbers come from molecular clouds relatively close to the solar neighborhood.  As a consistency check, we will also convert the Plume et al. $^{13}$CO J= 2$\\to$1 data to an H$_2$ column density more indirectly, using the  $^{12}$CO J=$2\\to1$ to $^{12}$CO J=$1\\to0$ and $^{12}$CO to $^{13}$CO J=$2\\to1$ line ratios discussed in the last section and the Galactic CO to H$_2$ conversion factor.  The uncertainties involved in this indirect determination will be larger than would be the case simply using the Bachiller \\& Cernicharo  density ratio, but it will provide a rough consistency check.  The average N(C)/N(H$_2$) ratio for the Plume et al. molecular clouds is 2.4~$\\times~10^{-5}$, a factor of four greater than we find for NGC~604-2. From I$_{[CI]}$/I$_{CO}$ = 0.17, derived in the previous section, N(C)/N(H$_2$) is 9~$\\times~10^{-6}$, nearly three times lower lower than what we found using the Bachiller \\& Cernicharo density ratio, and more consistent with our number for NGC~604-2.  However, as we mentioned, this second method is more indirect, assuming that $^{12}$CO J=$2\\to1$ to $^{12}$CO J=$1\\to0$ ratio and the $^{12}$CO to $^{13}$CO J=$2\\to1$ ratio from small Milky Way molecular clouds is applicable to the much larger NGC~604-2, which is located in a giant HII region.  It is reassuring that there is an agreement of a factor of three between these two calculations.  The most obvious differences between the Plume et al. clouds and NGC~604-2 are that NGC~604-2 is in a giant HII region, rather than a small HII region, and that NGC~604-2 is more metal poor than the HII region of Plume et al. Stronger dissociating radiation could lead to enhanced C creation through photodissociation.  It is also possible that the lower metallicity of NGC~604-2 plays a role; with less C, and therefore less CO, the self shielding of CO against dissociating radiation will be lower, allowing a greater fraction of the existing CO to be dissociated.  Of course, lower metallicity also means that the overall abundance of carbon will be lower, so it is not clear to what degree the enhanced photodissociation counters the lower metal abundance.  There is a difference in the physical size of the clouds -- the Plume et al. clouds are $\\sim$ 10 pc in size, compared to $\\sim$ 30 pc for NGC~604-2.  Perhaps differing N(C)/N(H$_2$) ratios represent a difference in the distribution of gas in the clouds, with the smaller galactic clouds having lower column density, and hence allowing greater penetration by dissociating radiation.  The models of Flower et al. give N(C)/N(H$_2$) ranging from $\\sim$ 10$^{-4}$ to 10$^{-7}$.  They can easily obtain values consistent with our observations, but the models have at least four free parameters, including the abundances of carbon and oxygen, the visual extinction, A$_v$, and degree of ionization.  With observations of just [CI] and CO we cannot distinguish between their various models.  The models of Spaans \\& van Dishoeck are more constrained because they attempt to model a specific Galactic molecular cloud, S140.  Their models assume a clumpy ISM which permits UV radiation to penetrate deep into the cloud.  They obtain column density ratios approaching ours for models with low volume filling factors (10\\% to 30\\%) and with relatively large clumps (0.4 to 0.6 pc).  Of course, S140 is different from NGC\\,604-2; it is illuminated by a single B0V star, not a giant HII region like NGC\\,604, so the agreement between the model and our observations should be taken with a grain of salt. However, lacking a detailed model specifically for NGC\\,604-2, a clumpy ISM would explain our observations.  \\subsection{[CI] Observations in Other Local Group Galaxies}  Stark et al. (1997) have measured the [CI] emission towards two regions within the Large Magellanic Cloud, N 159-W and 30 Dor N.  In the latter region, they obtain a I$_{[CI]}$/I$_{CO}$ line ratio of 0.26, twice the average for the Milky Way, but similar to what we have seen in NGC\\,604. They attribute this large ratio to the lower metallicity of the LMC. Bolatto et al. (2000, in press) find a similarly high line ratio, 0.23 $\\pm$ 0.03, in the low metallicity Local Group dwarf irregular galaxy IC~10. Our observations are fairly consistent with this picture; the metallicity of the NGC\\,604 region is nearly half solar (Vilchez et al. 1988) while the LMC is about one-third solar (8.37; Garnett 1990), and IC~10 is one-fourth solar metallicity (Lequeux et al. 1979).  \\placefigure{fig5}  Figure~5 shows a plot of I$_{[CI]}$/I$_{CO}$ versus metallicity for molecular clouds in Local Group galaxies. The Milky Way Galaxy is also included, with the data on I$_{[CI]}$/I$_{CO}$ coming from the COBE data of Wright et al. (1991). Because COBE did not detect CO J=1$\\to$0, we scaled the I$_{[CI]}$/I$_{CO J=2\\to1}$ ratio by a typical $^{12}$CO J=2$\\to$1 / $^{12}$CO J=1$\\to$0 line ratio of 0.7 (Sakamoto et al. 1994).  Although only eight points are plotted, the molecular clouds found near HII regions have consistently higher line ratios than those with no nearby HII regions. This is to be expected if radiation from HII regions is primarily responsible for producing atomic carbon by photodissociation of CO.  There also appears to be a trend among the clouds near HII regions for I$_{[CI]}$/I$_{CO}$ to decrease with increasing metallicity.  This trend might be expected if the column density of CO increases with metallicity, thereby more effectively shielding the interiors of the molecular clouds from dissociating radiation. It could also be a result of the metallicity dependence of the CO to H$_2$ conversion factor -- if I(CI)/N(H$_2$) remains constant, but I(CO)/N(H$_2$) decreases with metallicity, a similar effect might be observed.  More data are needed to confirm these trends.  \\subsection{The Origin of the [CI] Emission}  Figure~2 shows that the spatial distribution of [CI] emission is offset to the northwest relative to the CO.  This offset is not due to the pointing problems which were discussed and corrected in Section~2. The center of the giant HII region is northwest of NGC\\,604-2. This is where most of the O stars are found (Drissen et al. 1993) and hence the source of UV photons.  Thus the offset of the [CI] emission is towards the direction from which the photodissociating photons are incident, suggesting an edge-on morphology for the PDR in NGC\\,604-2.  The fact that [CI] emission is present into the center of the cloud indicates a non-negligible ionization fraction in the cloud's interior. At low ionizations the gas phase chemistry is dominated by reactions with H$^+_3$, which lead to the conversion of carbon into CO (Graedel et al. 1982).  One possibility to explain the necessary ionization inside the cloud is penetration by cosmic rays (e.g. Flower et al. 1994).  We argue that this mechanism is not likely to be an important one in the case of NGC\\,604-2 based upon the asymmetric spatial distribution of the [CI] emission.  For energies of 10$^{18}$ eV or less, a galactic magnetic field prevents cosmic rays from escaping a galaxy freely (Longair 1994).  These cosmic rays spiral around the magnetic field lines, lose their original direction of motion, and become isotropic.  Thus ionization from cosmic rays would not be expected to produce the asymmetry in [CI] emission that we see. In addition, Yang et al. (1996) have identified supernovae driven shells in NGC\\,604, and the spatial distribution of these shells (and hence the supernovae which could provide the cosmic rays) is not consistent with the [CI] distribution.  We cannot rule out a contribution to the creation of C by cosmic rays, but the lack of [CI] emission at the (0,-5) position, where CO emission is only 25\\% less than the cloud center (see Figure~1), suggests that any such contribution is small.  The (0,5) position lies on the opposite side of the cloud, but facing the center of the giant HII region, and I$_{[CI]}$ here is the strongest we have measured, stronger than at the center of the cloud, although the CO emission is nearly the same as at (0,-5).  This result suggests that the atomic carbon is being created by the photodissociation of CO by photons from the massive stars.  The penetration of the photons would then be permitted by a clumpy ISM with a moderate or low filling factor (e.g. Stutzki et al. 1988).  These photons penetrate partly through the cloud, even to the center, but do not penetrate all the way through, which explains the lack of [CI] emission on the far side.  Churchwell \\& Goss (1999) argue that the molecular cloud lies behind the HII region, based upon a measurement of extinction towards the cloud that is too low for the molecular mass of the cloud.  Our observations do not conflict with this idea, as long as NGC\\,604-2 is behind the bulk of the ionized gas along the line of sight, but still close enough to the HII region to receive the influx of UV photons that are responsible for dissociating CO and thus producing atomic carbon.  Also, if the cloud were too far behind the HII region, then a position like (0.-5) would not be shadowed by the cloud center, and thus would show [CI] emission."
    },
    "0002/hep-ph0002126.txt": {
        "abstract": "The evidence for the existence of dark matter in the universe is reviewed. A general picture emerges, where both baryonic and non-baryonic dark matter is needed to explain current observations. In particular, a wealth of observational information points to the existence of a non-baryonic component, contributing between around 20 and 40 percent of the critical mass density needed to make the universe geometrically flat on large scales. In addition, an even larger contribution from vacuum energy (or cosmological constant) is indicated by recent observations. To the theoretically favoured particle candidates for non-baryonic dark matter belong axions, supersymmetric particles, and of less importance, massive neutrinos. The theoretical foundation and experimental situation for each of these is reviewed. Direct and  indirect methods for detection of supersymmetric dark matter are described in some detail. Present experiments are just reaching the required sensitivity to discover or rule out some of these candidates, and major improvements are planned over the coming years. ",
        "introduction": "During the past few years, remarkable progress has been made in cosmology, both observational and theoretical. One of the outcomes of these rapid developments is the increased confidence that most of the energy density of the observable universe is of an unusual form, i.e., not made up of the ordinary matter (baryons and electrons) that we see around us in our everyday world.  The only realistic model for cosmology, and one that has become more solidly established by recent observations, is the big bang model. According to this Standard Model of cosmology, the universe has been and still is expanding from a compressed and hot phase, which existed some 15 billion years ago. The observational cornerstones of the big bang model -- the expansion of the universe, the fossil record of light elements sythesized during the first few minutes, and the existence still today of an intense thermal radiation field, the cosmic microwave background -- created when neutral atoms formed around 300 000 years after the big bang -- are more solid than ever. In fact, cosmology is entering an era of precision measurements, when information from all these three important processes is becoming accurate enough that it can be used to clarify the detailed structure and evolution of our Universe.  For example, big bang nucleosynthesis can be used to determine the baryon fraction of the matter density in the universe quite accurately \\cite{bbnreview}, and combined with analyses of galaxy cluster dynamics, supernova data and the cosmic microwave background radiation this gives  convincing arguments for the existence of a large amount of non-luminous, i.e., dark, matter. The matter content of the universe seems to be at least a factor of 5 higher than the maximum amount of baryonic matter implied by big bang nucleosynthesis. This dark matter is thus highly likely to be ``exotic'', i.e, non-baryonic \\cite{triangle}.  There are also indications, although still somewhat preliminary, of the existence of vacuum energy, corresponding to the famous ``cosmological constant'' that Einstein introduced but later rejected (although without very good reasons) in his theory of general relativity. An alternative possibility is a slowly varying scalar field, ``quintessence'' \\cite{steinhardt} which, like vacuum energy,  does not cluster at all (or on very large scales only). This  possibility has recently been given increased attention due to results from Type Ia supernova surveys \\cite{ariel,kirshner2} as will be discussed later in this review.  Although the existence of non-baryonic dark matter is now generally accepted by most of the astrophysical community, the nature of the dark matt\\-er is one of the outstanding questions in \\mbox{standard} cosmology \\cite{BGbook}.  In fact, since 1998 there is for the first time strong evidence for the existence of non-baryonic dark matter in the universe, in the form of massive neutrinos. This is due to the discovery of atmospheric neutrino oscillations in the Super-Kamiokande experiment \\cite{totsuka}. However, the natural neutrino mass scale of around 0.1 eV which is implied by the neutrino oscillation data is not large enough to influence cosmology in a dramatic way (although this type of non-baryonic dark matter would contribute about as much to the total mass density of the universe as do the visible stars).  Even if the mass of one of the neutrinos is higher (up to, say, 5 eV which would be possible if neutrinos are nearly degenerate in mass or if there exist a fourth, sterile neutrino), the mass density would not be large enough to explain the matter fraction of the cosmic average density. There are also arguments from galactic structure against an all-neutrino dark matter population. Therefore, it is natural to ask which other fundamental particles could be good dark matter candidates.  This review is organized as follows. In Section~\\ref{sec:basics} the general framework of modern cosmology is briefly reviewed. In Section~\\ref{sec:relic} the fundamental mechanism of thermal production of relic particles in the early universe is explained. In Section~\\ref{sec:amount} various estimation methods for the amount of different forms of mass and energy density on several length scales in the universe are discussed. In Section~\\ref{sec:baryons} the possible contribution from hidden ordinary matter (baryons) is treated, and the evidence for a population of dark baryonic objects, ``MACHO''s (MAssive Compact Halo Objects), displayed. The distribution of dark matter in galaxies, in particular the in Milky Way which most of the detection experiments are sensitive to, is discussed in Section~\\ref{sec:distribution}. Particle physics models for non-baryonic dark matter are treated in Section~\\ref{sec:models}, and the presently most discussed models for Weakly Interacting Massive Particles (``WIMP''s), namely the supersymmetric models, are reviewed in Section~\\ref{sec:susy}. These models, which have been developed in detail in the current literature, also serve as useful templates for other, yet unthought of, particle models for dark matter. Therefore, Section~\\ref{sec:detection} deals mainly with the methods of detection of the lightest electrically neutral supersymmetric particle,  the  neutralino, which is a quantum mechanical mixture of the supersymmetric partners of the photon, the $Z^0$ boson and the neutral scalar Higgs particles. Section~\\ref{sec:conclusions} ends the review with some concluding remarks and an outlook. ",
        "conclusions": "\\label{sec:conclusions} To conclude, we have seen that the existence of dark matter is more needed than ever, in order to explain a wealth of new observations. In addition to a large density of matter of unknown composition, the universe seems to contain also a large quantity of dark energy, of even more mysterious origin. Since the success of big bang nucleosynthesis combined with the measured intensity of the microwave background greatly constrains the amount of baryonic matter allowed, the major part of the matter density seems to be of non-baryonic origin.  The standard big bang model contains the required mechanism to create a relic density of electrically neutral, stable particles which, if their interactions are of electroweak strength, may be perfect candidates for dark matter. Foremost of these candidates are the hypothetical supersymmetric partners of electrically neutral ordinary particles, motivated by current thinking in particle physics. Also the more weakly interacting low-mass axions which were invented to solve the strong CP problem are attractive dark matter candidates, probed by sensitive on-going experiments.    Both  direct and indirect detection methods have the potential of investigating supersymmetric and similar massive particle  dark matter candidates. In particular, a combination of accelerator searches, new solid state, liquid or gas detectors, space-borne gamma-ray and air Cherenkov telescopes as well as neutrino telescopes may have the sensitivity needed to rule out or confirm the supersymmetry solution of the dark matter problem.  Since also the ex\\-perimental situa\\-tion con\\-cerning massive neu\\-trinos and axions is getting clearer, there is a chance to reach the goal of explaining the nature of the dark matter in the not too distant future.  However, rapid success is by no means guaranteed. It may be that the dark matter problem will plague the scientific community for a long time also in the new millennium. With the realization that dark vacuum energy may be an additional important component of our Universe, the mystery of its inner workings has deepened. However, this is a situation that is likely to inspire a young generation of physicists and astronomers to even more spectacular progress than the remarkable achievements since the 1930's when Zwicky first took notice of the dark matter problem. \\ack I wish to thank my collaborators, in particular Joakim Edsj\\\"o, Paolo Gondolo and Piero Ullio, for many helpful discussions. Also, I want to thank Elliott Bloom, Sandro Bottino, Jim Buckley, Per Carlson, Thibault Damour, Ulf Danielsson, Nicolao Fornengo, Tom Francke, Tom Gaisser, Ariel Goobar, Per Olof Hulth, Marc Kamionkowski, Lawrence Krauss, Joel Primack, Georg Raffelt, Hector Rubinstein, Pierre Sikivie, Joe Silk, and Michael Turner for useful input to this review.  This work has been supported in part by the Swedish Natural Science Research Council (NFR).  \\vskip .3cm"
    },
    "0002/astro-ph0002459_arXiv.txt": {
        "abstract": "We estimate the distribution of intrinsic shapes of APM galaxy clusters from the distribution of their apparent shapes. We measure the projected cluster ellipticities using two alternative methods. The first method is based on moments of the discrete galaxy distribution while the second is based on moments of the smoothed galaxy distribution. We study the performance of both methods using Monte Carlo cluster simulations covering the range of APM cluster distances and including a random distribution of background galaxies. We find that the first method suffers from severe systematic biases, whereas the second is more reliable. After excluding clusters dominated by substructure and quantifying the systematic biases in our estimated shape parameters, we recover a corrected distribution of projected ellipticities. We use the non-parametric kernel method to estimate the smooth apparent ellipticity distribution, and numerically invert a set of integral equations to recover the corresponding distribution of intrinsic ellipticities under the assumption that the clusters are either oblate or prolate spheroids. The prolate spheroidal model fits the APM cluster data best. ",
        "introduction": "Clusters of galaxies are the largest gravitationally collapsed objects in the universe and their internal dynamics and morphologies provide useful cosmological information. In recent years many studies of cluster shapes and orientations have showed that they are strongly elongated, maybe more so than elliptical galaxies, and they tend to point towards their neighbours (Carter \\& Metcalfe 1980; Bingelli 1982; Di Fazio \\& Flin 1988; Plionis, Barrow \\& Frenk 1991; De Theije, Katgert \\& van Kampen 1995). Plionis, Barrow and Frenk (1991) (hereafter PBF) have computed ellipticities and major axis orientations for the largest up to date sample of about 400 Abell clusters and found that their apparent shapes are consistent with those expected from a population of prolate spheroids. Support to the prolate spheroidal case was presented recently by Cooray (1999) analysing a sample of 25 Einstein X-ray clusters of Mohr et al. (1995). Struble and Ftaclas (1994) analysed a compilation of 344 Abell cluster ellipticities and found that rich clusters are intrinsically more spherical than poorer clusters. In the same framework McMilan et. al (1989) studied the ellipticities and orientations of 49 Abell clusters using Einstein X-ray data to trace the hot gas, and also found that the cluster potential is quite flat although less so than that found in optical studies. Buote \\& Canizares (1996) analyzed ROSAT PSPC images for 4 Abell clusters (including Coma) and, assuming hydrostatic equilibrium,they found ellipticities of order $\\epsilon_{mass}\\simeq 0.40-0.55$ (see also Canizares \\& Buote 1997).  Theoretical expectations regarding cluster shape and morphology have been investigated via N-body simulations, which show that the intrinsic shapes of simulated clusters are rather triaxial with an almost uniform distribution of shapes between prolate and oblate spheroids (cf. Frenk et al. 1988; Efstathiou et al. 1988). Detailed analysis of cluster morphological parameters and substructure, utilising the concept of power ratios (cf. Boute \\& Tsai 1994), can be used to constrain different cosmological models (cf. Thomas et al 1998; Valdarnini, Ghizzardi \\& Bonometto 1999).  It is obvious that information about the intrinsic shape of a cluster is lost when projected  on the plane of the sky. Many different studies have attempted to recover the distribution of intrinsic cluster shapes from the corresponding apparent distribution using inversion techniques based on the assumption that their orientations are random.  In this paper we use the APM catalogue (Dalton et al. 1997) to measure the projected shape distribution, corrected for various systematic effects, and hence attempt to estimate the intrinsic shape of clusters. The APM clusters are typically as rich as Abell $R=0$ clusters, but due to the careful identification procedure do not suffer from significant projection effects.  The plan of the paper is the following: In Section 2 we describe the APM galaxy and cluster survey, in Section 3 we present our projected cluster shape determination method and by using Monte Carlo simulations we establish its statistical robustness. We discuss how the foreground/background contamination (projection effects) affects the projected cluster shapes and present a statistical ellipticity correction procedure. The correction assumes that clusters are in dynamical equilibrium and therefore we exclude clusters with strong substructure. In Section 4 we invert the systematic bias-corrected projected ellipticity distribution to recover the intrinsic one. Our conclusions are presented in Section 5. ",
        "conclusions": "We have measured the projected ellipticities of all APM clusters using moments either of the individual galaxy distribution or of the smoothed galaxy distribution above some overdensity threshold. We have performed large sets of Monte Carlo simulations in order to test the statistical robustness of the two procedures, and conclude that the first method is strongly affected by the presence of background galaxies whereas the second method better recovers the underline true cluster projected ellipticity.  We devised a statistical Monte-Carlo procedure to correct the distribution of cluster ellipticities for the systematic errors introduced by discreteness effects, background galaxies and the method itself. The procedure involves estimating first the probability distribution of the true projected ellipticity for each APM cluster, and then random sampling this probability distribution to estimate the cluster's contribution to the corrected overall projected cluster ellipticity distribution. This method works well only for clusters that appear relaxed, with no significant substructure. `Therefore we have excluded, from our final cluster sample which contains 405 clusters, all the APM clusters that show evidence of significant substructure. Prior to the exclusion of these clusters the uncorrected axial ratio distribution of the whole APM cluster sample can be crudely approximated by a Gaussian with a mean of $\\simeq 0.65$ and a standard deviation of $\\simeq 0.15$. The corrected apparent axial-ratio distribution is significantly different showing a bump at $q\\simeq 0.46$ and having a significant contribution from apparently quasi-spherical systems.  Using the nonparametric kernel procedure we obtain a smooth estimate of the apparent APM cluster axial-ratio distribution. We assume that the APM clusters are a homogeneous population of either oblate or prolate spheroids and numerically invert the apparent distribution to obtain the intrinsic distribution.  The most acceptable model is provided by that of prolate spheroids. This result supports the view by which clusters form by accretion of smaller units along the large-scale structure (filament) in which they are embedded (cf. West 1994; West, Jones \\& Forman 1995). Such an accretion process would happen preferentially along the cluster major axis, which is typically aligned with the nearest cluster neighbour (cf. Bingelli 1982; Plionis 1994 and references therein).  Since cluster shapes and substructure are sensitive cosmological probes (Evrard et al. 1993), we plan to investigate these issues further using APM clusters and compare our results with theoretical expectations, provided by high-resolution N-body simulations of different cosmological models (cf. Thomas et al.1998)."
    },
    "0002/astro-ph0002490_arXiv.txt": {
        "abstract": "The Very Large Array (VLA) has been used in C-configuration to map an area $\\simeq0.3\\,{\\rm deg}^{2}$ at 1.4 GHz with $5\\sigma$ sensitivities of 0.305, 0.325, 0.380 and 0.450 mJy beam$^{-1}$ over four equal subareas. Radio properties are presented for 62 detected sources.  Deep optical imaging to Gunn $r\\simeq25$ mag using the Hale 5-m telescope covering $\\simeq0.21\\,{\\rm deg}^{2}$ is reported for a subset of 43 sources. This optical follow-up is much deeper than that of existing larger area radio surveys of similar radio sensitivity. Archival $J$, $H$ and $K$-band photometry from the Two-Micron All Sky Survey (2MASS) is also presented.  Using a robust likelihood-ratio technique, we optically identified 26 radio sources with probabilities $\\gtrsim80\\%$, nine are uncertain/ambiguous detections, and eight are empty fields. Comparisons with a stellar synthesis model that includes radio emission and dust reddening suggests that the near-infrared--to--optical emission in a small, bright sub-sample is reddened by `optically thin' dust with absorption $A_{V}\\simeq2-2.5$ mag, regardless of morphological type. This is consistent with other, more direct determinations of absorption. The radio--to--optical(--near-infrared) flux ratios of early-type galaxies require significant contamination in the radio by an active galactic nucleus (AGN), consistent with the current paradigm. Using our simple modeling approach, we also discuss a potential diagnostic for selecting Ultraluminous Infrared Galaxies (ULIGS) to $z\\simeq1.6$ from microJansky radio surveys. ",
        "introduction": "The advent of deep radio surveys reaching flux densities well below 1 mJy (Mitchell \\& Condon 1985; Windhorst et al. 1985; Oort 1987; Windhorst et al. 1993; Hopkins et al. 1998; Richards 2000) revealed a new population of faint sources more numerous than the AGN-powered radio galaxies dominating the strong-source population. This corresponds to a steepening of the differential source counts over non-evolving predictions at levels $\\lesssim4$ mJy. The faint counts suggest that significant evolution has occured over the redshift range spanned by the observed population. Photometric and spectroscopic studies (Thuan et al. 1984; Windhorst et al. 1985; Thuan \\& Condon 1987; Benn et al. 1993) suggest that the faint excess at 1.4 GHz is composed predominately of star-forming galaxies similar to the nearby starburst population dominating the {\\it Infrared Astronomical Satellite} ({\\it IRAS}) 60-$\\mu$m counts (Benn et al. 1993).  Indeed, this is supported by the strong correlation between radio (1.4 GHz) and far-infrared (60-$\\mu$m) flux densities of disk galaxies (Helou, Soifer \\& Rowan-Robinson 1985), implying a significant proportion of starburst galaxies at faint radio flux densities.  The overall observed source-count distribution from faint ($\\mu$Jy) to bright ($S_{1.4}\\gtrsim10$mJy) flux densities cannot be explained by starbursts alone. Evolutionary models of radio source counts need to invoke two separate populations (eg. Danese et al. 1987; Rowan-Robinson et al. 1993; Hopkins et al. 1998). Condon (1989) describes these populations as starbursts and monsters, each powered by different mechanisms: `starbursts' deriving their radio emission from supernova remnants and H{\\small II} regions and `monsters' from compact nuclear related activity (eg. active galactic nuclei; AGNs). The proportion of AGN is much greater at higher flux densities $S_{1.4}\\gtrsim10$ mJy (Kron, Koo \\& Windhorst 1985; Gruppioni et al. 1998), where a majority are associated with classical radio galaxies exhibiting extended (FRI and FRII-type) morphologies (Fanaroff \\& Riley 1974). The optical counterparts to sources at bright radio flux densities $S_{1.4}\\gtrsim1$ mJy is comprised mostly of ellipticals while at sub-mJy to $\\mu$Jy levels, the optical counterparts are identified as blue galaxies exhibiting peculiar (compact, interacting and irregular) morphologies (Kron, Koo \\& Windhorst 1985; Gruppioni et al. 1998). Studies of faint radio sources, namely their stellar population, how they evolve with redshift, and how they relate to local normal galaxies is progressing rapidly, however, much remains to be learned from the faintest ($\\mu$Jy) radio populations at redshifts $1\\lesssim z\\lesssim2$.  Radio surveys are insensitive to the effects of absorption by dust which is known to bias surveys severely at optical/UV wavelengths. This is particularly important for derivations of the cosmic history of star formation and its relation to hierarchical models of galaxy formation. Optical/UV studies have shown that there is an increase in both the space density of star-forming, morphologically disturbed galaxies (eg. Richards et al. 1999) and also, the global star formation rate with redshift to $z\\gtrsim1$ (eg. Madau et al. 1996). Similar evidence is emerging from studies of sub-mm sources (Blain et al. 1999) and amongst the faint radio population at $1\\lesssim z\\lesssim2$ - the redshift range probed by the deepest surveys (eg. Cram 1998; Haarsma et al. 1999).  There has been much speculation as to whether global star-formation rates (SFRs) derived from radio observations exceed those determined from optical/UV studies. Cram et al. (1998) note that systematic discrepancies may exist between the various star formation indicators, which are not well understood (see also Schaerer 1999). It is encouraging to see however that Haarsma et al. (2000) find global SFRs derived from 1.4GHz observations of the Hubble Deep Field to exceed optically detetermined values by a factor of a few out to  $z\\sim1$. Indeed, an analysis of Balmer decrements and optical--near-infrared colors in star forming galaxies by Georgakakis et al. (1999) from the Phoenix Deep Survey (Hopkins et al. 1998), finds evidence for visual extinctions from one to a few magnitudes. Currently, about 20\\% of existing micro-jansky radio samples remain unidentified to $I = 25$ mag in Hubble Deep Field images (eg. Richards et al. 1999). A majority of these could represent a significant population of dust-enshrouded starbursts and/or AGN at high redshift. These results are in support of efforts to further understand the dust properties of star-forming galaxies.  The primary aim of this paper is to introduce a new complete sample of radio sources selected at 1.4 GHz, uniformly selected over the flux range $S\\gtrsim0.3$ mJy (5$\\sigma$) from an area covering $\\simeq 0.3\\,{\\rm deg}^{2}$. Although much larger area surveys to deeper radio sensitivities have been carried out (eg. Hopkins et al. 1998), the present study reports the results of more sensitive optical observations. Archival near-infrared data for a subset of the sample are also presented. The near-infrared data is from the ongoing 2MASS project, and represents a unique aspect of this study in the identification of radio-selected starbursts. Although we currently lack valuable spectroscopic information, we combine radio--near-infrared--optical flux ratios, radio maps, and optical images to explore the properties of the entire sample. Our deep optical identifications provide the opportunity to asses the importance of dust in star-forming systems via the observed radio--to--optical and near-infrared--to--optical colors. Simple stellar synthesis models that include radio emission and reddening are used to constrain possible amounts of absorption.  This paper is organised as follows. In \\S 2, we discuss the radio observations and data reduction, present the radio catalog, and compare our results with data available from (shallower) all sky radio surveys. The optical photometric observations, their reduction and astrometric calibrations are discussed in \\S 3. Our method for radio-optical identification, the archival near-infrared data, radio-optical image overlays, and our optical/near-infrared catalog are presented in \\S 4. A study of the radio, near-infrared, and optical colors and constraints on synthesis models incorporating dust is presented in \\S 5. An application of our color-color analysis to select high-redshift ultraluminous infrared galaxies from deep radio surveys is discussed in \\S 6. All results and future prospects are summarised in \\S 7. ",
        "conclusions": "We have used the VLA radio telescope to image a contiguous $33\\times33\\,{\\rm arcmin}^{2}$ area to a (mean) limiting ($5\\sigma$) sensitivity of $\\simeq0.35$ mJy. From a total of 62 detections, the results of optical and near-infrared photometry are reported for 43 sources. Our optical photometry is more sensitive than previous optical follow-up studies of radio surveys of similar depth. Our main findings are:  \\indent (1) We have used a robust, likelihood-ratio method for determining optical identifications and their reliability. This method is seldom used in identification studies and is insensitive to assumptions concerning fluctuations in background source density and Gaussian error distributions. We assigned optical candidates to 26 radio sources with reliability $\\gtrsim80\\%$. Nine radio sources are uncertain and/or ambiguous, and eight are empty fields. Near-infrared photometry from the {\\it 2MASS} database was reported for 7 sources.  \\indent (2) The eight optical empty field sources all display compact and symmetric radio morphologies and most probably represent compact starbursts at $z\\lesssim0.3$ strongly obscured by dust. They may require at least 4 magnitudes optical extinction to account for their large radio--to--optical flux ratio compared to the identified population. Our conclusion for them being `compact starbursts' is very tentative as it is purely based on starburst versus AGN number statistics expected from sub-mJy radio surveys. Further deep infrared/optical imaging and spectroscopy will be necessary.  \\indent (3) Consistent with previous studies, our deep ($r\\simeq25$) optical imaging shows that the optical appearence can be divided into two classes according to radio flux-density: elliptical-like morphologies for $\\gtrsim2$ mJy, and peculiar or disturbed for $\\lesssim2$ mJy.  \\indent (4) Using a stellar synthesis model which includes radio emission and dust reddening, we find that the near-infrared--to--optical emission in a small, bright sub-sample is reddened by `optically thin' dust with $A_{V}\\simeq2-2.5$ mag, regardless of morphological type. This appears consistent with other more direct determinations. Consistent with previous studies, the radio emission from early-type systems seems to be powered by AGN rather than star-formation to account for their anomalously large radio--to--optical(--near-infrared) ratios.  \\indent (5) Our analysis shows that a radio/optical or radio/near-IR color selection technique could provide a potential means for detecting ULIG-type objects to $z\\sim1.6$.  Despite the lack of spectroscopic information, our study of a homogeneous population of faint radio sources has stressed the importance of dust on studies of intrinsic galaxy properties and their evolution at optical wavelengths. A future goal would be to obtain spectra, or multi-color optical/near-infrared photometry to better explore these sources and the validity of the simple stellar synthesis models presented in this paper. The ever improving resolution (and sensitivity) capabilities of optical/near-IR detectors over those feasible at (the longest) radio wavelengths requires robust identification techniques to better ascertain their properties. Likelihood-ratios provide one such technique. The present study complements other deep optical studies of faint radio sources to constitute a statistically significant sample for inferring their nature and importance to galaxy evolution."
    },
    "0002/astro-ph0002173_arXiv.txt": {
        "abstract": "Radial velocity measurements and simple sine-curve fits to the orbital velocity variations are presented for the third set of ten contact binary systems: CN~And, HV~Aqr, AO~Cam, YY~CrB, FU~Dra, RZ~Dra, UX~Eri, RT~LMi, V753~Mon, OU~Ser. All systems but two are contact, double-line spectroscopic binaries with four of them (YY~CrB, FU~Dra, V753~Mon, OU~Ser) being the recent discoveries of the Hipparcos satellite project. The most interesting object is V753~Mon with the mass-ratio closest to unity among all contact systems ($q = 0.970 \\pm 0.003$) and large total mass ($(M_1+M_2) sin^3i = 2.93 \\pm 0.06$). Several of the studied systems are prime candidates for combined light and radial-velocity synthesis solutions. ",
        "introduction": "\\label{sec1}  This paper is a continuation of the radial velocity studies of close binary stars, \\cite{LR99} (Paper I) and \\cite{RL99} (Paper II). The main goals and motivations are described in these papers. In short, we attempt to obtain modern radial velocity data for close binary systems which are accessible to the 1.8 meter class telescopes at medium spectral resolution of about R = 10,000 -- 15,000. Selection of the objects is quasi-random in the sense that we started with shortest period contact binaries. The intention is to publish the results in groups of ten systems as soon as reasonable orbital elements can be obtained from measurements evenly distributed in orbital phases. We are currently observing a few dozens of such systems. The rate of progress may slow down, as we move into systems with progressively longer periods.  This paper is structured in the same way as Papers~I and II in that it consists of two tables containing the radial velocity measurements (Table~\\ref{tab1}) and their sine-curve solutions (Table~\\ref{tab2}) and of brief summaries of previous studies for individual systems. The reader is referred to the previous papers for technical details of the program. In short, all observations described here were made with the 1.88 meter telescope of the David Dunlap Observatory (DDO) of the University of Toronto. The Cassegrain spectrograph giving the scale of 0.2 \\AA/pixel, or about 12 km/s/pixel, was used; the pixel size of the CCD was 19~$\\mu$m. A relatively wide spectrograph slit of 300~$\\mu$m corresponded to the angular size on the sky of 1.8 arcsec and the projected width of 4 pixels. The spectra were centered at 5185 \\AA\\ with the spectral coverage of 210 \\AA. The exposure times were typically 10 -- 15 minutes long.  The data in Table~\\ref{tab1} are organized in the same manner as in Paper~II. It provides information about the relation between the spectroscopically observed epoch of the primary-eclipse T$_0$ and the recent photometric determinations in the form of the (O--C) deviations for the number of elapsed periods E. It also contains, in the first column below the star name, our new spectral classifications of the program objects. The classification spectra, were obtained with a grating giving a dispersion of 0.62 \\AA/pixel in the range 3850 -- 4450 \\AA. The program-star spectra were ``interpolated'' between spectra of standard stars in terms of relative strengths of lines known as reliable classification criteria.  In the radial-velocity solutions of the orbits, the data have been assigned weights on the basis of  our ability to resolve the components and to fit independent Gaussians to each of the broadening-function peaks. Weight equal to zero in Table~\\ref{tab1} means that an observation was not used in our orbital solutions; however, these observations may be utilized in detailed modeling of  broadening functions, if such are undertaken for the available material. The full-weight points are marked in the figures by filled symbols while half-weight points are marked by open symbols. Phases of  the observations with zero weights are shown by short  markers in the lower parts of the figures; they were usually obtained close to the phases of orbital conjunctions.  All systems discussed in this paper but one have been observed for radial velocities for the first time. The only exception is RZ~Dra for which an SB1 orbit solution was obtained by \\cite{str46}. The solutions presented in Table~\\ref{tab2} for the four circular-orbit parameters, $\\gamma$, K$_1$, K$_2$ and T$_0$, have been obtained iteratively, with fixed values of the orbital period. First, two independent least-squares solutions for each star were made using the same programs as described in Papers~I amd II. Then, one combined solution for both amplitudes and the common $\\gamma$ was made with the fixed mean value of T$_0$. Next, differential corrections for $\\gamma$, K$_1$, K$_2$, and T$_0$ were determined, providing best values of the four parameters. These values are given in Table~\\ref{tab2}. The corrections to $\\gamma$, K$_1$, K$_2$, and T$_0$ were finally subject to a ``bootstrap'' process (several thousand solutions with randomly drawn data with repetitions) to provide the median values and ranges of the parameters. We have adopted them as measures of uncertainty of parameters in Table~\\ref{tab2}.  Throughout the paper, when the errors are not written otherwise, we express standard mean errors in terms of the last quoted digits, e.g. the number 0.349(29) should be interpreted as $0.349 \\pm 0.029$.  \\placetable{tab1}  \\placetable{tab2} ",
        "conclusions": "The paper brings radial velocity data for the third group of ten close binary systems that we observed at the David Dunlap Observatory. All, but RZ~Dra (which had SB1 radial velocity data), have never been observed spectroscopically; all ten are binaries with both components clearly detected so that they can be called SB2. All systems, but CN~And and RZ~Dra, which may be very close semi-detached systems, are contact binaries. We describe special features of the individual systems in the descriptions in Section~\\ref{sec2}. We note that again about half of the system are A-type contact binaries; the likely reasons why we prefer them over the W-type systems in our randomly drawn sample are given in the Conclusions to Paper~II. We also observed as EB2 systems two binaries, CN~And and RZ~Dra; they are most probably semi-detached binaries.  We do not give the calculated values of $(M_1+M_2) \\sin^3i = 1.0385 \\times 10^{-7}\\,(K_1+K_2)^3\\, P({\\rm day})\\,M_\\odot$ because in most cases the inclination angles are either unknown or not trustworthy. However, one case is very interesting here: The total mass of components of the system V753~Mon is very large, $M_1+M_2 > 2.93 M_\\odot$. Since {\\it too large velocity amplitudes\\/} are a rare phenomenon in the world of contact systems, this system requires special attention of the observers. The binary is also unique in having its mass-ratio exceptionally close to unity, $q_{sp}=0.970 \\pm 0.009$. Two other systems discovered by the Hipparcos mission are also very important and promise excellent combined solutions, YY~CrB and OU~Ser. It is important that both are high-velocity stars and have excellent parallaxes. And, finally, the recently discovered system HV~Aqr offers an excellent solution in view of the total, well-defined eclipses and very good radial velocity data."
    },
    "0002/astro-ph0002345_arXiv.txt": {
        "abstract": "We present for the first time a high-resolution X-ray spectrum of a Seyfert galaxy.  The Chandra-LETGS spectrum of NGC~5548 shows strong, narrow absorption lines from highly ionised species (the H-like and He-like ions of C, N, O, Ne, Na, Mg, Si, as well as \\ion{Fe}{xiv} -- \\ion{Fe}{xxi}).  The lines are blueshifted by a few hundred km/s.  The corresponding continuum absorption edges are weak or absent.  The absorbing medium can be modelled by an outflowing, thin and warm shell in photoionization equilibrium.  The absorption lines are similar to lower ionization absorption lines observed in the UV, although these UV lines originate from a different location or phase of the absorbing medium. Redshifted with respect to the absorption lines, emission from the \\ion{O}{viii} Ly$\\alpha$ line as well as the \\ion{O}{vii} triplet is visible.  The flux of these lines is consistent with emission from the absorbing medium.  The \\ion{O}{vii} triplet intensity ratios demonstrate that photoionization dominates and yield an upper limit to the electron density of $7\\times 10^{16}$~m$^{-3}$. ",
        "introduction": "Low to medium energy resolution X-ray spectra of AGN such as obtained by the Rosat or ASCA observatories showed the presence of warm absorbing material (see references in Kaastra \\cite{kaastra}).  This was deduced from broad band fits to the continuum, showing a flux deficit at wavelengths shorter than the expected edges of ions such as \\ion{O}{vii} and \\ion{O}{viii}.  The relation of this warm X-ray absorber to the medium that produces narrow UV absorption lines in \\ion{C}{iv} or \\ion{N}{v} is not clear, mainly due to a lack of sufficient constraints in the X-ray band. A major drawback of all previous X-ray studies of AGN has been the low spectral resolution, making it hard to disentangle any emission line features from the surrounding absorption edges, and prohibiting the measurements of Doppler shifts or broadening.  With the Chandra spectrometers it is now possible for the first time to obtain high-resolution X-ray spectra of AGN. ",
        "conclusions": "For the first time we see narrow absorption lines in the X-ray spectrum of an AGN.  However narrow absorption lines in Seyfert galaxies have been seen before in the UV band.  Shull and Sachs (\\cite{shull}) discovered narrow absorption features in the \\ion{C}{iv} and \\ion{N}{v} lines.  This was confirmed by Mathur et al.  (\\cite{mathur}) and studied in more detail by Mathur et al. (\\cite{mathur99}) and Crenshaw et al.  (\\cite{crenshaw}).  These last authors find at least 5 narrow absorption components in the \\ion{C}{iv}~1550~\\AA\\, and \\ion{N}{v}~1240~\\AA\\, lines.  These components are broadened by $\\sigma_v$=20--80~km/s, somewhat smaller than we find.  The rms width of the ensemble of UV absorption lines is 160~km/s for \\ion{C}{iv} and 260~km/s for \\ion{N}{v}, consistent with the effective line width of 270$\\pm$100~km/s that we find from our Chandra data.  Also, the average blueshift of the UV absorption lines ($-390$~km/s for \\ion{C}{iv} and $-490$~km/s for \\ion{N}{v}) is only slightly larger than what we find for the C, N and O lines ($-280\\pm 70$~km/s). Note that our wavelength scale has residual uncertainties of 100--300~km/s for most of our lines.  However, the column density of the lithium-like ions \\ion{C}{iv} and \\ion{N}{v} as derived by Crenshaw et al.  is 100 times smaller than the column density of the corresponding hydrogen-like ions that we find.  The difference may be attributed to either time variability (low column density during the UV observations), a high degree of ionization (hydrogenic ions dominating) or a stratified absorber (with UV and X-ray absorption lines originating from different zones).  We favour this last possibility.  This is supported by the fact that our simulations with XSTAR imply \\ion{C}{iv} and \\ion{N}{v} columns that are 100 times smaller than the measured values by Crenshaw et al.  Crenshaw \\& Kraemer (\\cite{crenshawk}) find that the weakest of the five dynamical components (their number 1) in the UV absorption lines has the highest outflow velocity ($-1056$~km/s).  Based upon the \\ion{N}{v} to \\ion{C}{iv} ratio, they argue that this component has the highest ionization parameter and could produce the oxygen continuum absorption edges as implied by the ASCA data. Our modelling with XSTAR also predicts column densities of \\ion{N}{v} and \\ion{C}{iv} close to the measured values for component 1.  But the outflow velocity of the X-ray absorber that we find is significantly smaller than the velocity of UV component 1.  However, Mathur et al.  (\\cite{mathur99}) identify component 3 ($-540$~km/s) as the most likely counterpart to the X-ray warm absorber.  We conclude that the detailed relation between UV and X-ray absorbers is still an open issue."
    },
    "0002/astro-ph0002035_arXiv.txt": {
        "abstract": "We show that, if many GRBs are indeed produced by the collapse of massive stars, GRBs and their afterglows provide a powerful probe of the very high redshift ($z \\gtrsim 5$) universe. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/astro-ph0002203_arXiv.txt": {
        "abstract": "The equation of state (EOS) of dense matter plays an important role in the supernova phenomenon, the structure of neutron stars, and in the mergers of compact objects (neutron stars and black holes).  During the collapse phase of a supernova, the EOS at subnuclear densities controls the collapse rate, the amount of deleptonization and thus the size of the collapsing core and the bounce density. Properties  of nuclear matter that are especially crucial are the symmetry energy and the nuclear specific heat.  The nuclear incompressibility, and the supernuclear EOS, play supporting roles.  In a similar way, although the maximum masses of neutron stars are entirely dependent upon the supernuclear EOS, other important structural aspects are more sensitive to the equation of state at nuclear densities.  The radii, moments of inertia, and the relative binding energies of neutron stars are, in particular, sensitive to the behavior of the nuclear symmetry energy. The dependence of the radius of a neutron star on its mass is shown to critically influence the outcome of the compact merger of two neutron stars or a neutron star with a small mass black hole.  This latter topic is especially relevant to this volume, since it stems from research prompted by the tutoring of David Schramm a quarter century ago. ",
        "introduction": "The equation of state (EOS) of dense matter plays an important role in the supernova phenomenon and in the structure and evolution of neutron stars.  Matter in the collapsing core of a massive star at the end of its life is compressed from white dwarf-like densities of about $10^6$ g cm$^{-3}$ to two or three times the nuclear saturation density, about $3\\cdot10^{14}$ g cm$^{-3}$ or $n_s=0.16$ baryons fm$^{-3}$. The central densities of neutron stars may range up to 5--10 $n_s$.  At densities around $n_s$ and below matter may be regarded as a mixture of neutrons, protons, electrons and positrons, neutrinos and antineutrinos, and photons.  At higher densities, additional constituents, such as hyperons, kaons, pions and quarks may be present, and there is no general consensus regarding the properties of such ultradense matter.  Fortunately for astrophysics, however, the supernova phenomenon and many aspects of neutron star structure may not depend upon ultradense matter, and this article will focus on the properties of matter at lower densities.  The main problem is to establish the state of the nucleons, which may be either bound in nuclei or be essentially free in continuum states. Neither temperatures nor densities are large enough to excite degrees of freedom, such as hyperons, mesons or quarks. Electrons are rather weakly interacting and may be treated as an ideal Fermi gas: at densities above $10^7$ g cm$^{-3}$, they are relativisitic.  Because of their even weaker interactions, photons and neutrinos (when they are confined in matter) may also be treated as ideal gases.  At low enough densities and temperatures, and provided the matter does not have too large a neutron excess, the relevant nuclei are stable in the laboratory, and experimental information may be used directly. The so-called Saha equation may be used to determine their relative abundances.  Under more extreme conditions, there are a number of important physical effects which must be taken into account.  At higher densities, or at moderate temperatures, the neutron chemical potential increases to the extent that the density of nucleons outside nuclei can become large.  It is then important to treat matter outside nuclei in a consistent fashion with that inside.  These nucleons will modify the nuclear surface, decreasing the surface tension.  At finite temperatures, nuclear excited states become populated, and these states can be included by treating nuclei as warm drops of nuclear matter.  At low temperatures, nucleons in nuclei are degenerate and Fermi-liquid theory is probably adequate for their description. However, near the critical temperature above which the dense phase of matter inside nuclei can no longer coexist with the lighter phase of matter outside nuclei, the equilibrium of the two phases of matter is crucial.  The fact that at subnuclear densities the spacing between nuclei may be of the same order of magnitude as the nuclear size itself will lead to substantial reductions in the nuclear Coulomb energy.  Although finite temperature ``plasma'' effects will modify this, the zero-temperature Wigner-Seitz approximation employed by Baym, Bethe \\& Pethick~\\cite{BBP} is usually adequate.  Near the nuclear saturation density, nuclear deformations must be dealt with, including the possibilities of ``pasta-like'' phases and matter turning ``inside-out'' ({\\it i.e.}, the dense nuclear matter envelopes a lighter, more neutron-rich, liquid).  Finally, the translational energy of the nuclei may be important under some conditions.  This energy is important in that it may substantially reduce the average size of the nuclear clusters.  An acceptable way of bridging the regions of low density and temperature, in which the nuclei can be described in terms of a simple mass formula, and high densities and/or high temperatures in which the matter is a uniform bulk fluid, is to use a compressible liquid droplet model for nuclei in which the drop maintains thermal, mechanical, and chemical equilibrium with its surroundings.  This allows us to address both the phase equilibrium of nuclear matter, which ultimately determines the densities and temperatures in which nuclei are permitted, and the effects of an external nucleon fluid on the properties of nuclei.  Such a model was originally developed by Lattimer {\\it et al.\\/} \\cite{LPRL} and modified by Lattimer \\& Swesty~\\cite{LS}.  This work was a direct result of David Schramm's legendary ability to mesh research activities of various groups, in this case to pursue the problem of neutron star decompression.  After the fact, the importance of this topic for supernovae became apparent. ",
        "conclusions": ""
    },
    "0002/astro-ph0002086_arXiv.txt": {
        "abstract": "A physical description of the fundamental plane of Galactic globular clusters is developed which explains all empirical trends and correlations in a large number of cluster observables and provides a small but complete set of truly independent constraints on theories of cluster formation and evolution in the Milky Way.  In a very good first approximation, the internal structures and dynamics of Galactic globulars are described by single-mass, isotropic \\markcite{kin66}King (1966) models and are thus characterized by four nominally independent physical parameters. Within this theoretical framework, it is shown that (1) 39 regular (non--core-collapsed) globulars with measured central velocity dispersions all share a common core mass-to-light ratio, $\\Upsilon_{V,0}=1.45\\,M_\\odot\\,L_\\odot^{-1}$, and (2) 109 regular globulars both with and without direct observations of $\\Upsilon_{V,0}$ show a very strong correlation between global binding energy and total luminosity, regulated by Galactocentric position: $E_b=7.2\\times10^{39}\\ {\\rm erg}\\ (L/L_\\odot)^{2.05}(r_{\\rm gc}/ 8\\,{\\rm kpc})^{-0.4}$. The observational scatter about either of these two constraints can be attributed fully to random measurement errors, making them the {\\it defining equations of a fundamental plane} (FP) to which real clusters are confined in the larger, four-dimensional parameter space of general King models. They are shown to underlie a pair of bivariate correlations first found, and used to argue for the existence of a globular cluster FP, by \\markcite{djo95}Djorgovski (1995). A third, weaker correlation, between clusters' total luminosities and King-model concentration parameters, is related to the (non-random) distribution of globulars {\\it on} the FP. With $L$, $\\Upsilon_{V,0}$, $E_b$, and the central concentration $c$ thus chosen as the four physical quantities that define any single globular cluster, the FP equations for $\\Upsilon_{V,0}$ and $E_b(L,r_{\\rm gc})$ are used to {\\it derive} expressions for {\\it any other} observable in terms of only $L$, $r_{\\rm gc}$, and $c$. Results are obtained for generic King models and applied specifically to the globular cluster system of the Milky Way. ",
        "introduction": "The properties of globular clusters in the Milky Way offer empirical constraints not only on their own formation and evolution, but also on the history of our Galaxy as a whole and, to a potentially large extent, on the formation of generic star clusters---a key element of the star formation process itself. The many observed correlations between various internal structural and dynamical parameters of the clusters (core and half-light radii, surface brightnesses, velocity dispersions, etc.) that have been found by several authors (e.g., \\markcite{bro84}Brosche \\& Lentes 1984; \\markcite{kor85}Kormendy 1985; \\markcite{vdb91}van den Bergh, Morbey, \\& Pazder 1991; \\markcite{vdb94}van den Bergh 1994; \\markcite{djo94}Djorgovski \\& Meylan 1994; \\markcite{djo95}Djorgovski 1995; \\markcite{bel96}Bellazzini et al.~1996; \\markcite{mur92}Murray \\& Lin 1992) constitute an important set of such constraints.  The apparently large number of these correlations (see especially \\markcite{djo94}Djorgovski \\& Meylan 1994) can be somewhat misleading, however, in that it is impossible for all of them to be independent. This follows from the fact that Galactic globulars are structurally and dynamically exceedingly simple, being very well described by single-mass, isotropic \\markcite{kin66}King (1966) models. Since these models are defined fully by the specification of just four independent physical parameters, it is not unreasonable to expect that the same should be true, at least in a first-order approximation, of real clusters; but there can then be at most three truly distinct correlations between their various intrinsic properties (although these may be supplemented, in general, by dependences on ``external'' factors such as Galactocentric position or cluster metallicity).  It is clearly of interest, then, to systematically reduce the globular cluster data in the Milky Way to the minimum number of independent physical relationships required for a complete description of the observed properties of the ensemble. The establishment of such a compact set of empirical constraints---which should serve better as a target for theories of cluster formation and evolution---is the main goal of this paper. To reach it, two important results will be used as starting points.  First, the two best globular cluster correlations known to date (i.e., the only two with scatter that can be fully accounted for by observational uncertainties) are the bivariate dependences presented by \\markcite{djo}Djorgovski (1995) for a subset (roughly one-third) of the Galactic globulars. One of these correlations involves only cluster core parameters, while the other includes properties measured at their half-light radii. \\markcite{djo95}Djorgovski found these correlations by a statistical principal-components analysis, and he offered no physical interpretation of them; but he pointed out that their very existence suggests that globular clusters may in practice be confined---much like elliptical galaxies and other large stellar systems are---to a ``fundamental plane'' in the higher-dimensional space of physical parameters available to them in principle.  Second, an almost completely unexamined (but, as it happens, very closely related) issue has to do with the dependence of binding energy, $E_b$, on mass (or luminosity) among globular clusters. This was studied some time ago by \\markcite{sai79}Saito (1979), who also investigated the question for dwarf and giant elliptical galaxies. He found that $E_b\\propto M^{1.5}$ or so---apparently consistent with very simple virial-theorem arguments---for globulars and giant ellipticals alike. However, his observational sample included only ten relatively bright globular clusters. The data required for a calculation of binding energy are now available for many more clusters than this and they are obviously of higher quality than those available to \\markcite{sai79}Saito. A total energy $E_b$ is the only obvious physical property that is {\\it not} routinely estimated from observations of Galactic globulars; but it is arguably important---if for no other reason than to complete a systemization of the data---to clarify the empirical behavior of this quantity.  In what follows, catalogued data are used to estimate binding energies within the context of single-mass, isotropic \\markcite{kin66}King (1966) models for 109 ``regular'' (non--core-collapsed) globular clusters in the Milky Way. The data and necessary aspects of the models are summarized in the next Section and in a pair of Appendices. Section 3 presents the binding energies themselves, and also derives core mass-to-light ratios, $\\Upsilon_{V,0}$, for 39 of these clusters with measured central velocity dispersions. It is shown that (1) $\\Upsilon_{V,0}$ is a constant for all clusters, and (2) $E_b$ correlates tightly with total cluster luminosity (or mass), scaling as $E_b\\propto L^{2.05\\pm0.1}$. This is significantly (if not surprisingly) different from \\markcite{sai}Saito's (1979) result, and perhaps also from the relation obeyed by elliptical galaxies (although this also appears not to have been reviewed in any detail since \\markcite{sai79}Saito's work).  Section 3 also establishes the choice of $L$, $E_b$, $\\Upsilon_{V,0}$, and the King-model concentration parameter $c\\equiv\\log\\,(r_t/r_0)$, as the four ``independent'' physical properties used to define any globular cluster. The tight constraints found on $\\Upsilon_{V,0}$ and $E_b(L)$ then imply that real clusters have (at most) only two free parameters, and they do indeed occupy a fundamental plane in the theoretical space of King models. The influence of environment on this plane, as reflected in cluster Galactocentric radii and metallicities, is also examined in \\S3; it is shown to be confined to a dependence on $r_{\\rm gc}$ in the {\\it normalization} of the $E_b$--$L$ scaling.  In \\S4, an orthogonal coordinate system called $\\epsilon$-space (analogous to the so-called $\\kappa$-space of elliptical galaxies; \\markcite{ben92}Bender, Burstein, \\& Faber 1992) is constructed to obtain a face-on view of the fundamental plane. This shows that globulars are not distributed uniformly over it; essentially, the parameter $c$, which describes the internal density profile of a cluster, is also correlated with total luminosity or mass (see also \\markcite{vdb94}van den Bergh 1994; \\markcite{bel96}Bellazzini et al.~1996). Alternate representations of the fundamental plane, in the form of a number of bivariate correlations between cluster core and half-light parameters, are also developed in \\S4. It is shown that the two correlations of \\markcite{djo95}Djorgovski (1995) are equivalent to the two constraints $\\Upsilon_{V,0}=constant$ and $E_b\\propto L^{2.05}$ and, thus, that \\markcite{djo95}Djorgovski identified the same fundamental plane that is developed here.  Section 5 and Appendix A then use the defining equations of the fundamental plane, along with manipulations based on generic properties of King model, to derive a set of simple relations that ultimately describe the variation of {\\it any} globular cluster observable as a function of $L$, $c$, and $r_{\\rm gc}$ {\\it only}. These may themselves be manipulated to explain any observed correlation between any other combination of cluster variables. Thus, the fundamental plane indeed reduces a large array of globular cluster data to a very small set of independent constraints for models of their formation and evolution. These are set out in \\S6, which summarizes the paper.  It should be noted that the choice of ``basic'' cluster properties adopted here---total luminosity (mass), binding energy, mass-to-light ratio, and concentration parameter---is not a unique possibility, nor is it necessary in any rigorous sense. It will be noted, for example, that the scaling of binding energy with luminosity is essentially equivalent to other, previously known correlations between half-light radii and luminosities ($R_h\\sim L^0$; \\markcite{vdb91}van den Bergh et al.~1991) and between core velocity dispersions and luminosities ($\\sigma_0\\sim L^{0.52}$; e.g., \\markcite{djo94}Djorgovski \\& Meylan 1994). However, this does not make a description in terms of $E_b$ any less correct or valid. In fact, this framework seems particularly amenable to theoretical studies of cluster evolution and formation both: Many simulations of the former follow the evolution of $c$ and $E_b$ over time, and a discussion of the $E_b(L,r_{\\rm gc})$ dependence in the latter context will be taken up in a subsequent paper (McLaughlin, in preparation). ",
        "conclusions": "If they are described by single-mass, isotropic \\markcite{kin66}King (1966) models, globular clusters are fully defined, in general, by specifying just four independent physical parameters. These were chosen in this paper to be a mass-to-light ratio ($\\log\\,\\Upsilon_{V,0}$), total binding energy ($\\log\\,E_b$), central concentration [$c=\\log\\,(r_t/r_0)$], and total luminosity ($\\log\\,L$). It has been shown that (1) all 39 regular (non--core-collapsed) Galactic globular clusters with measured core velocity dispersions have a common mass-to-light ratio (\\S3.1), $$\\log\\,(\\Upsilon_{V,0}/M_\\odot\\,L_\\odot^{-1})=0.16\\pm0.03\\ ;$$ and (2) if it is assumed that this also holds for the other regular clusters in the Milky Way with no velocity data, but for which measured King-model structural parameters are available, the full cluster system (109 objects in all, excluding core-collapsed members) shows a very well-defined dependence of binding energy on total luminosity, modulated by Galacocentric radius (\\S3.2): $$\\log\\,\\left(E_b/{\\rm erg}\\right)= \\left[\\left(39.86\\pm0.40\\right)-0.4\\,\\log\\,\\left(r_{\\rm gc}/8\\,{\\rm kpc} \\right)\\right]+\\left(2.05\\pm0.08\\right)\\,\\log\\,\\left(L/L_\\odot\\right)\\ .$$ The scatter about each of these relations is fully accounted for by the typical observational errorbars on $\\log\\,\\Upsilon_{V,0}$ ($\\simeq0.2$ dex) and $\\log\\,E_b$ ($\\simeq0.5$ dex).  With $\\Upsilon_{V,0}$ essentially a fixed constant and $E_b$ known precisely (at a given Galactocentric position) as a function of $L$, only two of the four basic cluster properties are truly independent; globular clusters in the Milky Way are confined to a narrow, two-dimensional subregion---a fundamental plane (FP)---in the larger, four-dimensional space of King models. The distribution of clusters on the plane is then determined essentially by their $c$- and $L$-values. They are not scattered randomly, however, as there also exists a correlation between $c$ and $\\log\\,L$. Still, this does {\\it not} necessarily mean that the locus of globular clusters is only one-dimensional: a linear $c$--$\\log\\,L$ relation is of poorer quality than either of the empirical constraints on $\\log\\,\\Upsilon_{V,0}$ and $\\log\\,E_b$, and although there {\\it might} exist some more complicated function which gives a comparably one-to-one dependence of concentration on luminosity, one has yet to be found (\\S\\S3 and 4).  Regardless, the mean trend of $c$ vs.~$\\log\\,L$ is independent of Galactocentric position and cluster metallicity, as is the value of $\\log\\,\\Upsilon_{V,0}$ and the slope of the correlation between $\\log\\,E_b$ and $\\log\\,L$ (\\S3.3). The normalization of the $E_b(L)$ scaling is also independent of [Fe/H], but it decreases towards larger Galactocentric radii. The equation above accounts for this, and the $r_{\\rm gc}$ dependence there describes the {\\it full extent of environmental influences} on the globular cluster FP.  It was shown in \\S4 that cluster concentration parameters (and, to a lesser extent, binding energies) correlate more tightly with luminosity for globulars at $r_{\\rm gc}>8$ kpc than for those at smaller radii. Thus, since dynamical evolution has likely been less effective at erasing initial conditions outside the Solar circle---and since simple arguments indicate that clusters could remain on or close to the FP as they evolve---it may well be that most properties of the fundamental plane were set more or less at the time of globular cluster formation. Quantitative calculations of cluster evolution over a Hubble time in the Galactic potential should be applied to check this.  A three-dimensional ``$\\epsilon$-space'' was constructed in \\S4.1 from a simple transformation of the cluster parameter space left after specifying $\\log\\,\\Upsilon_{V,0}=constant$ (and after removing the effects of Galactocentric radius) in order to obtain a directly face-one view of the FP. It could be interesting to attempt a similar construction for elliptical galaxies and clusters of galaxies, i.e., to interpret their fundamental planes in terms of binding energy as well, in order to make a more direct comparison with the globular cluster FP.  The equations for $\\Upsilon_{V,0}$ and for $E_b$ as a function of $L$, which define the FP, were shown in \\S4.2 to be equivalent to the two strong bivariate cluster correlations used by \\markcite{djo95}Djorgovski (1995) to argue for the existence of a fundamental plane for globular clusters. \\markcite{djo95}Djorgovski (and \\markcite{bel98}Bellazzini 1998) offered a simple virial-theorem argument (which was shown here to be incomplete) as an explanation for one of these correlations, but did not interpret the other. The results of this paper have put his results on a firm physical footing.  Finally, since $\\Upsilon_{V,0}$, $c$, $E_b$, and $L$ completely define globular clusters, any correlations connecting any of their other properties, or any other trends with Galactocentric radius, necessarily derive from generic properties of \\markcite{kin66}King (1966) models combined with the empirical FP relations and their environmental dependence. Appendix A outlined such derivations quite generally, and specialized results for the Milky Way cluster ensemble were presented in \\S5.  The picture that has been developed here is a simplification in that it works specifically in the context of single-mass and isotropic models for globular clusters. This approximation is evidently an excellent one, but because of it the results of this paper can say nothing about the stellar mass functions or possible velocity anisotropy in observed clusters. Nor have the kinematics of clusters within the Milky Way (e.g., orbital eccentricities, or bulk rotation of the metal-rich subsystem) been considered in any way. These issues aside, however, the preceding discussion may be distilled into a small set of four main facts to be explained by theories of cluster formation and evolution:  \\begin{itemize}  \\item All cluster properties are independent of metallicity.  \\item The core mass-to-light ratio, $\\Upsilon_{V,0}=(1.45\\pm0.1)\\,M_\\odot\\, L_\\odot$, does not vary significantly with cluster luminosity or Galactocentric position.  \\item The concentration parameter, which controls the shape of a cluster's internal density profile, correlates with luminosity. The dependence of $c$ on $\\log\\,L$ is much more significant outside the Solar circle than inside, but even there the relationship is not obviously one-to-one. Its basic form is, however, independent of $r_{\\rm gc}$.  \\item Binding energy is intrinsically a function of luminosity and is regulated by Galactocentric position: $E_b\\propto L^{2.05}r_{\\rm gc}^{-0.4}$. Since there is no evidence for significant variations in the {\\it global} mass-to-light ratios of clusters in the Milky Way, $E_b$ is inferred to scale with total mass in the same way as with luminosity.  \\end{itemize}  Once again, the balance of current evidence suggests that these characteristics of the fundamental plane---and, thus, most of the systematics of the Galactic globular cluster system---were essentially fixed by the cluster formation process. Moreover, it is not implausible that this process was controlled largely by a single intrinsic protocluster parameter---the initial gas mass---and adjusted by an external influence depending on Galactocentric radius. Further discussion along these lines is left for future work (McLaughlin, in preparation). Regardless of any interpretive details, however, it clearly will be important to determine the extent to which the globular cluster systems of other galaxies can be described in the same simple terms that apply in the Milky Way."
    },
    "0002/astro-ph0002279_arXiv.txt": {
        "abstract": "We present the results of a follow--up {\\it ASCA} observation and multicolour optical and near-infrared photometry carried out at the 3.5-m Italian National Telescope Galileo of SAXJ~1353.9$+$1820. This object, serendipitously discovered by {\\it BeppoSAX} in the 5--10 keV band, has been spectroscopically identified as a red quasar at $z$=0.217. The combined X--ray and optical--infrared data reveal the presence of a moderately luminous X--ray source ($\\sim$ 10$^{44}$ erg s$^{-1}$) obscured by a column density of the order of 10$^{22}$ cm$^{-2}$ in a otherwise optically passive early-type galaxy. The implications for the nature of red quasars and their possible contribution to the hard X--ray background are briefly outlined. ",
        "introduction": "The {\\it BeppoSAX} High Energy Large Area Survey (HELLAS) has discovered a large population of hard X--ray sources, which account for about 20--30 per cent of the cosmic X--ray background (XRB) in the 5--10 keV energy range (Fiore et al. 1999). This band provides an efficient tool to discriminate between accretion-powered sources, like the Active Galactic Nuclei (AGN) are generally thought to be, and sources dominated by the starlight component. Moreover, hard X--ray selection is less affected by obscuration with respect to other bands, making possible the study of the nuclear continuum without any relevant contamination by reprocessed radiation. Indeed the optical identification process indicates that the large majority of the HELLAS sources are AGN (La Franca et al., in preparation). Among these, SAXJ~1353.9$+$1820 can be considered one of the most intriguing sources. After its discovery in the context of the HELLAS survey, it was spectroscopically identified as a radio-quiet AGN at a redshift of 0.2166 (Fiore et al. 1999).  Its optical spectrum is dominated by starlight, as the presence of H and K plus Mg~I absorption lines clearly indicates. The H$\\alpha$ equivalent width (about 90 $\\AA$), the absence of the H$\\beta$ and the {\\it BeppoSAX} hardness ratio are suggestive of significant reddening of the nuclear radiation. The hard X--ray luminosity, L$_{5-10~keV}$ $\\sim$ 1.4 $\\times$ 10$^{44}$ erg s$^{-1}$, the presence of a broad Balmer line and the optical spectrum properties allow to classify this object as a fairly low-luminosity, red quasar. \\\\ This class of objects, originally discovered in radio--selected samples (Smith \\& Spinrad 1980), is characterized by red optical colours ($B-K$ up to 8, Webster et al. 1995). On the basis of {\\it ROSAT} PSPC observations and optical spectroscopy Kim \\& Elvis (1999) pointed out that a significant fraction (from a few percent up to 20 per cent) of soft X--ray selected radio--quiet quasars belong to this class. The origin of the observed `redness' may be ascribed to dust absorption, to intrinsic red colours ($\\ga$ 70 per cent of the red sources of the Parkes sample has $\\la$ 30 per cent of contribution to B--K by their host galaxies, Masci et al. 1998) or to an excess of light in K band rather than a dust-induced deficit in B (Benn et al. 1998). Whatever is the origin of the red colours, it is likely that a large fraction of quasars could have been missed by the usual selection techniques in the optical band (Webster et al. 1995). If this is the case, red quasars could constitute a sizeable fraction of the absorbed AGN population needed to explain the hard X--ray background spectrum (i.e. Comastri et al. 1995) especially if the optical reddening is associated with X--ray absorption. Hard X--ray observations provide the most efficient way to select these objects; indeed already two candidates have been found among the first optical identifications of HELLAS sources (Fiore et al. 1999). In order to better understand the spectral properties of these objects we have started a program of multiwavelength follow--up observations of a sub--sample of HELLAS sources. Here we present the first results obtained in the X--ray band with {\\it ASCA} and at optical--infrared wavelengths with the 3.5-m Italian National Telescope Galileo (TNG) at La Palma (OIG and ARNICA photometric cameras). Throughout the paper $H_{0}$ = 50 km s$^{-1}$ Mpc$^{-1}$ and $q{_0}$ = 0 are assumed. ",
        "conclusions": "The {\\it ASCA} observation confirms the presence of a bright and moderately absorbed ($N_{\\rm H}$ $\\sim$ 6 $\\times$ 10$^{21}$ cm$^{-2}$) nucleus with a flat hard X--ray spectrum, in agreement with the BeppoSAX hardness ratio analysis. Assuming for the underlying continuum slope the average value of quasars in the same energy range, the absorption column density could be as high as $N_{\\rm H}$ $\\sim$ 3 $\\times$ 10$^{22}$ cm$^{-2}$ if some 20 per cent of the nuclear radiation is not absorbed. Even if the X--ray data alone does not allow to distinguish between the two possibilities, we can safely conclude that SAXJ~1353.9$+$1820 harbours a mildly obscured, luminous (L$_{\\rm 2-10 keV}$ $\\simeq$ 1.3 $\\times$ 10$^{44}$ erg s$^{-1}$) active nucleus.  Not surprisingly high energy observations of quasars characterized by similar dust reddened optical continua do reveal the presence of absorption by cold (IRAS 23060+0505, Brandt et al. 1997a) and/or warm (IRAS 13349+2438, Brandt et al. 1997b) gas.  A basic step forward in understanding the nature of this red quasar and of red quasars as a whole, is provided by optical plus near--IR studies (see Maiolino et al., in preparation, for further details). The surface brightness profiles are consistent with those of an elliptical galaxy, and the optical colours (U$-$B=0.43, B$-$V=1.40, V$-$R=0.83 and R$-$I=0.70) agree with the properties of a early-type galaxy at $z$=0.2 (Fukugita et al. 1995).  In order to obtain a self--consistent description of the optical-IR properties we fitted the photometric points with a two-components model consisting of an old stellar population template (10$^{10}$ yr, Bruzual \\& Charlot 1993) and a moderately absorbed ($A_{\\rm V} \\simeq$ 2 mag, corresponding to $N_{\\rm H} \\simeq$ 4 $\\times$ 10$^{21}$ cm$^{-2}$ for Galactic dust--to--gas ratio, {\\rm i.e.} Bohlin, Savage \\& Drake 1978) quasar spectrum template (Elvis et al. 1994; Francis et al. 1991). As shown in Fig.~4, the combination of these two spectra provides a good description of the observed data (possibly with the exception of the J photometry, which deviates by 1.4 sigma). The quality of the fit is acceptable ($\\chi^2_{\\rm r}$ = 1.3 when both the uncertainties in the photometric data and in the template spectra are taken into account) indicating that most of the optical and near--IR flux is dominated by star--light.  \\begin{figure} \\psfig{figure=h8final.ps,width=0.48\\textwidth,angle=0} \\caption[]{SAXJ~1353.9$+$1820 photometric (optical $+$ NIR) points fitted with a synthetic model of an evolved early-type galaxy (dashed line) plus the contribution of a moderately absorbed quasar (dotted line). The sum is represented by the solid line.} \\label{fig4} \\end{figure}  The present results add further evidence on the hypothesis of substantial gas and dust absorption as an explanation of the observed properties of this red quasar. Even more interesting is that the active nucleus peers only at X--ray energies and possibly at wavelenghts longward of 2 $\\mu$m (see Fig.~4), while at optical wavelengths SAXJ~1353.9$+$1820 looks like a normal evolved elliptical galaxy. If this behaviour applies also to other objects it may well be that a significant fraction of obscured AGNs resides in otherwise normal passive galaxies. These nuclei would have been completely missed in optical quasar surveys because of their extended morphology and galaxy-like colours. If the column density is of the order of  10$^{22}$ cm$^{-2}$ or higher, their fraction could be understimated also in soft X--ray surveys. If this is the case, the fraction of red objects among radio quiet quasars ($\\sim$ 3--20 per cent, Kim \\& Elvis 1999) should be considered as a lower limit. This may have strong implications for the XRB synthesis models, which in their simplest version (i.e. Madau, Ghisellini \\& Fabian 1994; Comastri et al. 1995) predict a large number of high-luminosity absorbed quasars (called type 2 QSO). Despite intensive optical searches, these objects appears to be elusive indicating a much lower space density than that of lower luminosity Seyfert 2 galaxies (Halpern, Eracleous \\& Forster 1998) and calling for a substantial revision of AGN synthesis models for the XRB (Gilli, Risaliti \\& Salvati 1999). \\\\  An alternative possibility (see Comastri 2000) is that X--ray obscured AGN show a large variety of optical properties including those of SAXJ~1353.9$+$1820. It is worth noting that column densities as high as 10$^{23.5}$ cm$^{-2}$ have been detected in Broad Absorption Line QSO (Gallagher et al. 1999), in some UV bright soft X--ray weak QSO (Brandt et al. 1999), and in a few HELLAS sources optically identified with broad line blue quasars (Fiore et al. 1999). It is thus possible that the sources responsible for a large fraction of the XRB energy density are characterized by a large spread in their optical to X--ray properties. The $\\alpha_{ox}$ spectral index, defined as the slope joining the 2500 $\\AA$ and the 2 keV flux densities, is usually employed to measure the optical to X--ray ratio. Not surprisingly, absorbed objects are characterized by values of $\\alpha_{ox}$ ($>$ 1.8) much steeper than the average value of bright unabsorbed quasars and Seyfert galaxies, $\\alpha_{ox}$ $\\simeq$ 1.5 (Laor et al. 1997; Yuan et al. 1998), while the faint nuclear UV flux density and the relatively bright 2 keV flux of SAXJ~1353.9$+$1820 correspond to $\\alpha_{ox}$ $\\simeq$ 1, which is quite flat but not unusual for red quasars (Kim \\& Elvis 1999).  A detailed discussion on the quasars contribution to the XRB is beyond the purposes of this Letter. Here we note that as far as the X--ray spectral properties and the bolometric luminosity of $\\simeq$ 10$^{45}$ erg s$^{-1}$ (estimated using the average SED of Elvis et al. 1994 with the measured $\\alpha_{ox}$) are concerned, SAXJ~1353$+$1820 can be classified as a high--luminosity absorbed AGN.  Future sensitive X--ray observations with Chandra and XMM coupled with optical spectropolarimetry data would be extremely helpful to better understand the nature of red quasars and to estimate their contribution to the XRB."
    },
    "0002/astro-ph0002423_arXiv.txt": {
        "abstract": "{\\it BV\\/} photographic photometries of 789 stars in the region of the open clusters NGC1750 and NGC1758 are derived from a set of 8 photometric plates. According to the astronomical data including proper motions, positions and membership probabilities of individual stars, several astrophysical parameters for these two clusters are determined, such as HR diagrams, ages, distances, luminosity functions, masses and kinematics, etc. It is found that the  distances are ($525\\pm 48$pc)  and ($794\\pm 73$pc) with ages of $1.5\\times10^8$yr and $6.3\\times10^8$yr for NGC1750 and NGC1758 respectively. Both clusters show no significant mass segregation effects because of their relatively young ages. The analysis of the proper motions imply that their velocity distributions are isotropy.\\  \\ Keywords{ open cluster: individual -- NGC1750, NGC1758 -- photometry -- parameters} ",
        "introduction": "Open clusters, as systems of stars having a common origin, provide a very powerful tool on studying stellar evolution history.  The homogeneity of photometric characteristics of stars in a cluster and their dynamics indicate that the cluster stars should have formed from one and the same primordial cloud within a relatively short time scale. Therefore, almost all member stars in a cluster should have roughly same age and chemical composition.  Furthermore, open clusters can be used to understand both the formation and kinematics of the Galactic disk due to their wide age and mass distributions.  For these reasons, open clusters constitute one of the most important  research fields in observational and theoretical astronomy.  Ages, distances, masses, luminosity functions and mass functions of open clusters are their basic researches.  The region of the open cluster NGC1750 has been paid more and more attention mainly because it is a complex area including two open clusters (NGC1750 and NGC1758) and the Taurus dark clouds, which located in the anticenter direction of the Galaxy. Cuffey(1937) was the first one to study this area systematically, obtaining extensive photographic photometry of stars in blue and red photometric bands to a limiting magnitude about $R\\sim 14\\,$mag. However he did not clearly point out the existence of two open clusters in this region. He considered the whole area as NGC1746 with distance about 590pc and \\bmr$\\,\\approx\\,$0.30mag.  More recently, Galad\\'{i}-Enr\\'{i}quez et al. (1998a, hereafter GJTR; 1998b; 1998c) did a series of studies of the open clusters NGC1750 and NGC1758, in which {\\it UBVRI\\/} CCD photometry of 3224 stars within the 45\\min$\\,\\times\\,$45\\min\\ area was presented. At the same time, they combined different plates from several observatories to obtain proper motions for 45036 stars, and {\\it BVR\\/} photographic photometry of 39762 stars within 2\\Deg3$\\,\\times\\,$2\\Deg3 in this area  was completed down to $V\\,\\sim\\,18.5\\,$mag. Several physical parameters for NGC1750 and NGC1758 were discussed, including their positions, sizes, density profiles, extinctions, distances, ages, luminosity functions and masses, etc. Tian et al.\\ (1998, hereafter TZSS) obtained high-precision proper motions and membership probabilities for 540 stars within 1\\Deg5$\\,\\times\\,$1\\Deg5 in this area using 20 plates taken over a period of up to 68 years from Z\\v{o}-S\\'{e} 40cm Astrograph in Shanghai Astronomical Observatory, Chinese Academy of Sciences. In that work the two open clusters, NGC1750 and NGC1758, are successfully separated from each other. The core radii of the clusters have been estimated to be 17\\Min20 and 2\\Min25 respectively.  In this paper, we will present the results of photographic photometry of 789 stars to a limit of $V\\sim16\\,$ mag within a 35\\min$\\,\\times\\,$35\\min\\ area round the center of NGC1750. Combining the results of proper motions and membership probabilities for individual stars, we will more deeply study some astrophysical parameters of the two open clusters, including their H-R diagrams, distances, ages, masses, mass functions and kinematics. The program is as follows. In sec. 2, we will introduce the data reduction and compare with previous work. Detailed discussion on the basic astrophysical researches are presented in Sec.3. Summary is arranged in the final section. ",
        "conclusions": "In the present paper, based on the proper motions, photometries and membership probabilities of individual stars in the region of NGC1750 (TZSS, GJTR), we investigate basic astrophysical properties for two dynamically independent open clusters, NGC1750 and NGC1758. After detailed discussions on the photometric data and membership probabilities, the analysis samples with star number of 311 and 23 for NGC1750 and NGC1758 are constructed respectively. Comparing ZAMS (Mermilliod 1981; Schaller et al. 1992), we obtain the distances for these two clusters of $(525\\pm 48)\\,$pc for NGC1750 and $(794\\pm 73)\\,$pc for NGC1758 with their extinction being considered. Furthermore, many methods for the age determination are adopted to estimate the average ages of $(2.2\\pm0.6)\\times10^8 \\rm yr$ and $(7.8\\pm1.6)\\times10^8 \\rm yr$ with the observed masses 390$\\,$M$_\\odot$ and 40$\\,$M$_\\odot$ for NGC1750 and NGc1758 respectively.  According to  the results of membership determination, luminosity functions and mass functions are given at the same time. It can be concluded that there exist no significant mass segregation effects for both clusters, which is consistent with  the fact that their dynamical relaxation have not undergone thoroughly.  Finally, the velocity distributions of member stars for these two clusters are also discussed. It is found that both clusters seem to be isotropy in velocity space. Moreover, it is worth nothing that our statistical results could not be enough certain for NGC1758 because of its small number of member stars observed."
    },
    "0002/astro-ph0002109_arXiv.txt": {
        "abstract": "The {\\sc Aladin} interactive sky atlas, developed at CDS, is a service providing simultaneous access to digitized images of the sky, astronomical catalogues, and databases. The driving motivation is to facilitate direct, visual comparison of observational data at any wavelength with images of the optical sky, and with reference catalogues.  The set of available sky images consists of the STScI Digitized Sky Surveys, completed with high resolution images of crowded regions scanned at the MAMA facility in Paris.  A Java WWW interface to the system is available at:  {http://aladin.u-strasbg.fr/}. ",
        "introduction": "\\subsection{The CDS} The Centre de Donn\\'ees astronomiques de Strasbourg (CDS) defines, develops, and maintains services to help astronomers find the information they need from the very rapidly increasing wealth of astronomical information, particularly on-line information.  In modern astronomy,  cross-matching data acquired at different wavelengths is often the key  to the understanding of astronomical phenomena, which means that astronomers have to use data and information produced in fields in which they are not specialists. The development of tools for cross-identification of objects is of particular importance in this context of multi-wavelength astronomy.  A detailed description of the CDS on-line services can be found, e.g., in Egret et al.\\ (\\cite{cds-amp2}) and in Genova et al.\\ (\\cite{cds-hub}, \\cite{cds}, \\cite{cds2000}), or at the CDS web site\\footnote{\\emph{Internet address:} http://cdsweb.u-strasbg.fr/}.   \\subsection{The ALADIN Project}  Several sites currently provide on-line access to digitized sky surveys at different wavelengths: this is, for instance, the case of Digitized Sky Survey (DSS) at STScI (Morrison \\cite{1995adass...4..179M}), and of similar implementations at other sites, providing quick access to cutouts of the compressed DSS images. {\\sc SkyView} at HEASARC (McGlynn et al. \\cite{skyview}) can generate images of any portion of the sky at wavelengths in all regimes from radio to gamma-ray. Some of these services  provide simultaneous access to images and to catalogue data. The {\\sc SkyCat} tool, recently developed at ESO (Albrecht et al.\\ \\cite{skycat}), addressed this concern in the context of the European Southern Observatory scientific environment (in view of supporting future users of the Very Large Telescope); {\\sc SkyCat} uses a standardized syntax to access heterogeneous astronomical data sources on the network.   {\\sc Aladin} has been developed independently by the CDS since 1993 as a dedicated tool for identification of astronomical sources -- a tool that can fully benefit from the whole environment of CDS databases and services, and that is designed in view of being released as a multi-purpose service to the general astronomical community.  {\\sc Aladin} is an interactive sky atlas, allowing the user to visualize a part of the sky, extracted from a database of images from digitized surveys and observational archives, and to overlay objects from the CDS catalogues and tables, and from reference databases ({\\sc Simbad} and NED), upon the digitized image of the sky.  It is intended to become a major cross-identification tool, since it allows recognition of astronomical sources on the images at optical wavelength, and at other wavelengths through the catalogue data. Expected usage scenarios include multi-spectral approaches such as searching for counterparts of sources detected at various  wavelengths, and applications related to careful identification of astronomical objects. {\\sc Aladin} is also heavily used for the CDS needs of catalogue and database quality control.  In the case of extensive undertakings (such as checking the astrometric quality for a whole catalogue), it is expected that {\\sc Aladin} will be useful for understanding the characteristics of the catalogue or survey, and for setting up the parameters to be adjusted while fine tuning the cross-matching or classification algorithms, by studying a sample section of objects or fields.  A discussion of the usage of such a tool for cross-identification can be found in Bartlett \\& Egret (\\cite{xid-179}), where it is shown how \\emph{training sets} are used to build likelihood ratio  tests.  The {\\sc Aladin} interactive atlas is available in three modes: a simple previewer, a Java interface, and an X-Window interface.  We describe here mostly the Java interface which is publicly accessible on the World-Wide Web. ",
        "conclusions": ""
    },
    "0002/astro-ph0002446_arXiv.txt": {
        "abstract": "While in the hierarchical model of structure formation, groups of galaxies are believed to be the scaled-down version of clusters of galaxies, a similarity breaking in the fundamental laws may occur on the  group scale, reflecting a transition between galaxy-dominated and intracluster medium dominated properties. In this paper, we present an extensive study of the relations between the X-ray luminosity ($L_x$), the temperature ($T$) of hot diffuse gas and the velocity dispersion ($\\sigma$) of galaxies for groups and clusters of galaxies, based on the largest sample of 66 groups and 274 clusters drawn from literature. Our best fit $L_x$-$T$ and $L_x$-$\\sigma$ relations for groups read $L_x\\propto T^{5.57\\pm1.79}\\propto \\sigma^{2.35\\pm0.21}$, which deviates remarkably from those for clusters: $L_x\\propto T^{2.79\\pm0.08}\\propto \\sigma^{5.30\\pm0.21}$. The significance of these correlations have been justified by both the co-consistency test and the Kendall's $\\tau$ statistics. We have thus confirmed the existence of similarity breaking in the $L_x$-$T$ and $L_x$-$\\sigma$ relations between groups and clusters as claimed in previous work, although the best fit $\\sigma$-$T$ relations remain roughly the same in both systems: $\\sigma\\propto T^{0.64}$. Alternatively, the significant disagreement between the observationally fitted $L_x$-$T$ and $L_x$-$\\sigma$ relations for groups and those expected from a perfect hydrostatic equilibrium hypothesis indicates that the X-ray emission of individual galaxies and the non-gravitational heating must play a potentially important role in the dynamical evolution of groups. Therefore, reasonable caution should be exercised in the cosmological applications of the dynamical properties of groups. ",
        "introduction": "It has been well established that there exists a strong correlation between the X-ray determined bolometric luminosity $L_x$, the X-ray temperature $T$ of the intracluster gas and the optical measured velocity dispersion $\\sigma$ of the cluster galaxies (Wu, Xue \\& Fang 1999; hereafter Paper I and references therein). A precise determination of these correlations is important not only for the study of dynamical properties and evolution of clusters themselves but also for distinguishing different cosmological models including the prevailing estimate of the cosmic density parameter $\\Omega_M$ through a combination of the baryon fraction  of clusters and the Big Bang Nucleosynthesis (White et al. 1993; David, Jones \\& Forman 1995). For instance, if the observed $L_x$-$T$ relation follows $L_x\\propto T^{2}$ (e.g. Quintana \\& Melnick 1982; Markevitch et al. 1998), we would arrive at the conclusion that the X-ray emission is primarily due to purely gravitational heating and thermal bremsstrahlung, and the baryon fraction $f_b$  of clusters can be representative of the Universe, i.e., $f_b$ provides  a robust estimate of $\\Omega_M$ (see Paper I). However, if the observed $L_x$-$T$ relation appears to be $L_x\\propto T^{3}$ (e.g. White, Jones \\& Forman 1997), other non-gravitational heating mechanisms must be invoked in order to give gas sufficient excess energy (e.g. Ponman, Cannon \\& Navarro 1999; Wu, Fabian \\& Nulsen 1999; Loewenstein 1999), unless the baryon fraction of clusters is requited to vary with temperature (David et al. 1993). The latter challenges the standard model of structure formation.   Another critical issue is whether the $L_x$-$T$, $L_x$-$\\sigma$ and $\\sigma$-$T$ relations for clusters of galaxies can naturally extend to group scales. In the hierarchical model of structure formation, groups are believed to be the scaled-down version of clusters, and the underlying gravitational potentials of groups and clusters are similar when scaled to their virial radii (e.g. Navarro, Frenk \\& White 1995). It is expected that groups and clusters should exhibit similar correlations between $L_x$, $T$ and $\\sigma$, provided that gas and galaxies are in hydrostatic equilibrium with the underlying gravitational potential of groups and their X-ray emissions are produced by thermal bremsstrahlung. Indeed, in the present Universe  a majority of the baryons may be bound in the gravitational wells of groups  (Fukugita, Hogan \\& Peebles 1998). All groups may contain hot X-ray emitting gas with $kT$ around or less than $1$ keV (e.g. Price et al. 1991; Ponman et al. 1996), and most of them should be detectable with future sensitive observations. Without the knowledge about the dynamical properties of groups characterized by the $L_x$-$T$, $L_x$-$\\sigma$ and $\\sigma$-$T$ relations, it could be misleading if the presently estimated baryon fraction of groups is immediately used as a cosmological indicator. Meanwhile, any difference in these correlations between clusters and groups will be helpful for our understanding of the formation and evolution of structures on scales of $1$--$10$ Mpc and of the significance of non-gravitational heating mechanisms.   The pioneering work of constructing  the  $L_x$-$\\sigma$ relation for groups was carried out by Dell'Antonio, Geller \\& Fabricant (1994). They found a significant flattening in the relation for groups with $\\sigma$ below $300$ km s$^{-1}$, as compared with the same relation for rich clusters. An extensive study of the $L_x$-$T$, $L_x$-$\\sigma$ and $\\sigma$-$T$ relations for groups was soon made by Ponman et al.(1996), based on 22 Hickson's compact groups (HCG) whose X-ray emissions are detected. The most remarkable result is the steepening ($L_x\\propto T^{8.2}$) of the $L_x$-$T$ relation, in contrast with the X-ray properties of clusters ($L_x\\propto T^{2.5}$--$T^{3}$), while the significant flattening in the  $L_x$-$\\sigma$ relations for groups claimed by Dell'Antonio et al. (1994) was only marginally detected. Using the RASSCALS group catalog, Mahdavi et al. (1997, 2000) obtained a much shallower $L_x$-$\\sigma$ relation ($L_x\\propto \\sigma^{0.37\\pm0.3}$) for groups with $\\sigma<340$ km s$^{-1}$ than the above results. Interestingly, Mulchaey \\& Zabludoff (1998) studied 12 poor ROSAT groups but arrived at an opposite conclusion that the X-ray component in groups follows the same $L_x$-$T$-$\\sigma$ relations as those in clusters. Namely, groups are indeed low-mass versions of clusters.   Theoretically, several investigations have been made on the possible reasons why the $L_x$-$T$ relation flattens with the increase of scale and temperature (e.g. Cavaliere, Menzi \\& Tozzi 1997, 1999; Valageas \\& Schaeffer 1999). In particular, if the $L_x$-$T$ and $L_x$-$\\sigma$ relations show a significant departure from the expectations of $L_x\\propto T^{2}\\propto \\sigma^{4}$ within the framework of thermal bremsstrahlung and hydrostatic equilibrium, the previous estimate of the amount of baryonic matter in groups and its application to the determination of cosmic density parameter would become questionable. Consequently, one may need to study whether the observed X-ray emission of groups has partially arisen from the individual galaxies ( Dell'Antonio et al. 1994) or other non-gravitational mechanisms such as those suggested by the preheated ICM model (e.g. Ponman et al. 1999) and the shock model (Cavaliere et al. 1997).   In this paper, we would like to update and then compare the $L_x$-$T$, $L_x$-$\\sigma$ and $\\sigma$-$T$ relations for groups and clusters, using all the available X-ray/optical data in literature especially the new measurements of poor clusters and groups. We wish to achieve a better statistical significance, which may essentially allow us to closely examine the question as to whether the similarity of these relations would break down on group scales. Throughout the paper we assume a Hubble constant of $H_0=50$ km s$^{-1}$ Mpc$^{-1}$ and a flat cosmological model of $\\Omega_0=1$. ",
        "conclusions": "\\subsection{Global properties}   Groups and clusters constitute large and nearly virialized dynamical systems in the Universe.  While the distribution of dark matter  in these systems may exhibit a regularity such as the universal density profile, the hot intragroup/intracluster gas could be affected by non-gravitational heating mechanisms (e.g. star formation), especially for poor systems like groups of galaxies where even individual galaxies may make a nonnegligible contribution to the X-ray emission (Dell'Antonio et al.  1994; Mulchaey \\& Zabludoff 1998). Therefore, the standard scenario that groups and clusters should be the similar dynamical systems at different mass ranges is only applicable to the distribution and evolution of the dark matter particles. Whether or not the hot gas can be used as a good tracer of  the underlying gravitational potential wells depends on how significant the X-ray emission of individual galaxies and the non-gravitational heating would be, which in turn depends on how massive a dynamical system will be. Because groups of galaxies connect individual galaxies to clusters of galaxies, one may expect to detect the transition between ``galaxy-dominated'' and ``intracluster medium dominated'' properties occurring on group scales (e.g. Dell'Antonio et al.  1994).   Such a scenario has essentially been justified by our analysis of the correlations between X-ray luminosity, temperature and velocity dispersion for  groups and clusters, although the current data especially for groups still have rather large uncertainties. It is remarkable that we have detected the similarity breaking in the $L_x$-$T$ and $L_x$-$\\sigma$ relations between massive systems (clusters) and low mass ones (groups), indicating that the overall dominant X-ray emission and heating mechanisms are quite different in these two systems. Nevertheless, for the $\\sigma$-$T$ relation we have not found an apparent discrepancy between the two systems. Our result essentially agrees with the previous findings of Dell'Antonio et al. (1994), Ponman et al. (1996) and Mahdavi et al. (1997, 2000).    \\subsection{Uncertainties}   Major uncertainties in the presently fitted $L_x$-$T$, $L_x$-$\\sigma$ and $\\sigma$-$T$ relations for groups and clusters come from the sparse (X-ray) data sets of poor clusters and groups. Although an observational determination of these relations does not in principle require the completeness of the sample, the reliability and significance of our fitting can be seriously affected by the small number statistics. The marginal co-consistency between the three relations for groups may be partially due to the small group sample. It is generally believed that poor clusters and groups should all contain  hot X-ray emitting gas (e.g. Price et al. 1991; Ponman et al. 1996). However, most of them can be hardly detected by current X-ray observations because of their low temperatures $T\\sim0.1$--$1$ keV. Therefore, our group sample suffers from a selection bias, in which the present X-ray data are acquired by different authors with different criteria. For instance, about 1/3 of our group sample are HCGs, and the steepening of the $L_x$-$T$ relation for groups may have partially arisen from the contribution of these HCG populations.   Our statistical results are sensitive to the linear fitting methods, especially for the $L_x$-$T$ and $L_x$-$\\sigma$ relations. This arises because the observed quantities, $L_x$, $T$ and $\\sigma$, all have measurement uncertainties, while OLS method ignores scatters around the horizontal variable (e.g. $T$ or $\\sigma$). At this point, the ODR fitting technique is strongly recommended.  However, not all the data points have information about their measurement uncertainties, or some error bars are difficult to evaluate. In this case, we have simplified the problem and employed the average values instead, in order to maximally use the available data points, which may have yielded further uncertainties in the ODR fitted relations. Note that the measurement uncertainties  in $L_x$ are relatively smaller than those in $T$ and $\\sigma$ (see Fig.1 and Fig.2). This explains the significant difference in the resultant $L_x$-$T$ and $L_x$-$\\sigma$ relations between  OLS and ODR methods.   \\subsection{Physical implications}  Under the assumption that both gas and galaxies are in hydrostatic equilibrium with the underlying gravitational potential of group/cluster, we would expect that the total X-ray luminosity via thermal bremsstrahlung scales as (e.g. Paper I) \\begin{equation} L_x\\propto T^{2.5}f_b^2 r_{c}^{-1} S_{gas}, \\end{equation} and \\begin{equation} L_x\\propto \\sigma^{4} T^{1/2} f_b^2 r_{c}^{-1} S_{gal}, \\end{equation} where $f_b$ is the (gas) baryon fraction, $r_c$ is the core radius and $S$ is the so-called dimensionless structure factor which depends strongly on the power-index of gas/galaxy profile (e.g. the conventional $\\beta$ model or the King model) but weakly on the core radius. Using the emission weighted temperature instead of $T$ in eq.(1) only leads to a modification to $S_{gas}$. Additionally, the velocity dispersion of galaxies and the temperature of hot X-ray emitting gas satisfy \\begin{equation} \\sigma \\propto  T^{1/2}. \\end{equation}   From our fitted $\\sigma$-$T$ relations alone,  groups are still consistent with $\\sigma \\propto  T^{1/2}$ within $2\\sigma$ uncertainties, while a significant deviation from what is expected under the scenario of isothermal and hydrostatic equilibrium is found for clusters. The latter is consistent with a number of similar studies on the issue (e.g. White, Jones \\& Forman 1997; Wu et al. 1998; 1999). Yet, it is unlikely that the intragroup gas is in the state of a perfect isothermal and hydrostatic equilibrium with the underlying gravitational potential as allowed by the $\\sigma$-$T$ relation within its $95\\%$ significance interval. This point can be further demonstrated by the apparent disagreement between the theoretical prediction of eq.(1), $L_x\\propto T^2$, and our fitted $L_x$-$T$ relation  for groups in Table 2, $L_x\\propto T^{5.57\\pm1.79}$, unless the baryon fraction is assumed to vary with gas temperature. Meanwhile, the best fit $L_x$-$\\sigma$ relation for groups, $L_x\\propto \\sigma^{2.35\\pm0.21}$, differs from the theoretical expectation (eq.[2]) in conjunction with eq.(3): $L_x\\propto \\sigma^{4}$.  Taking these results as a whole, we feel that the currently available optical/X-ray data have already provided convincing evidence for either the failure of isothermal and hydrostatic equilibrium hypothesis for intragroup gas or the significant contribution of X-ray emission from the individual group galaxies. Therefore, any cosmological applications of these fitted relations without considering these effects could be misleading.   The physical implications of the $L_x$-$T$, $L_x$-$\\sigma$ and $\\sigma$-$T$ relations for clusters have been extensively studied in Paper I. For the group sample, there exists an apparent difference between the theoretically predicted $L_x$-$T$ and $L_x$-$\\sigma$ relations (eq.[1]) and the observationally determined ones. Dell'Antonio et al. (1994) attributed the flattening of $L_x$-$\\sigma$ relation to the additional X-ray emission of individual group galaxies, while Ponman et al (1996) interpreted the steepening of the  $L_x$-$T$ relation as the result of the wind injection from galaxies. The recent detection of the X-ray wakes in pool cluster A160 may give a strong support to these scenarios, i.e., a large fraction of X-ray emission in groups and poor clusters can be associated with individual galaxies (Drake et al. 2000). It thus deserves to  explore whether this scenario can  quantitatively account for the reported $L_x$-$T$-$\\sigma$ relations for groups. Alternatively, other mechanisms such as preheating by supernovae and substructure merging may also contribute extra heating to the intragroup medium.   On the other hand, the prevailing determination of the baryon fraction of groups via hydrostatic equilibrium and thermal bremsstrahlung could be seriously contaminated by the X-ray emission of individual group galaxies and other non-gravitational heating. This might account for the relatively large variations of the presently derived baryon fractions  among different groups. It is unlikely that a robust estimate of the baryon fractions of groups and poor clusters within the framework of conventional dynamical models and their cosmic evolution can be achieved without a better understanding of the various heating mechanisms. Indeed, we cannot exclude the possibility that some puzzles about the properties of the baryon fractions of groups and clusters are due to the contamination of the X-ray emission of individual galaxies and the non-gravitational heating. For instance, the standard model predicts an increase in the baryon fraction with radius, and a universal value of $f_b$ at large radii cannot be reached unless the X-ray temperature is required to rise in some clusters (Wu \\& Xue 2000 and references therein). The presence of such puzzle has essentially prevented the baryon fractions of groups and clusters from the cosmological applications. Inclusion of the contributions of other emission/heating sources in the estimate of the total gravitational masses of groups and clusters may lead to a re-arrangement of the radial distribution of baryon fraction.  Further work will thus be needed to explore whether these effects can resolve the above puzzle."
    },
    "0002/astro-ph0002393_arXiv.txt": {
        "abstract": "\\noindent Gigahertz Peaked Spectrum (GPS) sources form a key element in the study of the onset and evolution of radio-loud AGN, since they are most likely the young counterparts of extended radio sources. Here we discuss space-VLBI observations of GPS sources, which enable us to obtain unprecedented angular resolution at frequencies near their spectral turnovers. Observed peak brightness temperatures of $10^{10.5-11}$ Kelvin indicate that synchrotron self absorption is responsible for their spectral turnovers. This is in close agreement with previous size $-$ spectral turnover statistics for GPS sources. The combination of these new space-VLBI observations with ground-based VLBI observations taken at an earlier epoch, confirm the young ages for the most compact GPS galaxies of several hundred years. ",
        "introduction": "Although radio-loud Active Galactic Nuclei (AGN) have been studied for several decades, still not much is known about their birth and subsequent evolution. The recent identification of a class of very young radio sources can be considered as a major breakthrough in this respect, since it has opened many unique opportunities for radio source evolution studies.  Unfortunately, the nomenclature and use of acronyms in this field of research is rather confusing. This is mainly caused by the different ways in which young radio sources are selected. Selection of young sources is made in two ways, the first based on their broadband radio spectra, and the second based on their compact morphology. A convex shaped spectrum, peaking at about 1 GHz distinguishes young radio sources from other classes of compact radio sources. In this case they are called Gigahertz Peaked Spectrum ({\\bf GPS}) radio sources (eg. O'Dea etal. 1991, O'Dea 1998). Similar objects, which are typically an order of magnitude larger in size, have their spectral turnovers shifted to the $10 - 100$ MHz regime, causing them to be dominated  at cm wavelengths by the optically thin parts of their spectra. These are called Compact Steep Spectrum ({\\bf CSS}) radio sources to distinguish them from the general population of extended steep spectrum sources (eg. Fanti et al. 1991).  On the other hand, young radio sources are found in multi-frequency VLBI surveys, in which they can be recognised by  compact jet/lobe-like structures on both sides of their central core. They are called Compact Symmetric Objects ({\\bf CSO}, Wilkinson et al. 1994). Their double sided structures clearly distinguish them from the large majority of compact sources showing one-sided core-jet morphologies. This implies that the luminosities of CSO are unlikely to be substantially enhanced by Doppler boosting. Larger versions of CSOs are subsequently called Medium Symmetric Objects ({\\bf MSO}) and Large Symmetric Objects ({\\bf LSO}).  The overlap between the classes of CSO and GPS galaxies is large and we believe that they can be considered to be identical objects. However, note that a substantial fraction of GPS sources are optically identified with high redshift quasars, which in general show core-jet structures (Stanghellini et al. 1997). The relationship between GPS quasars and GPS galaxies/CSO is not clear and under debate (Snellen et al. 1999). We therefore believe it is wise to restrict evolution studies to GPS galaxies and CSOs.  \\subsection*{Evidence for youth}  Although it was always speculated that GPS sources were young objects, only recently has strong evidence been found to support this hypothesis. Monitoring several GPS sources over a decade or more using VLBI, allowed Owsianik \\& Conway (1998) and Owsianik, Conway \\& Polatidis (1998) to measure the hotspot advance speeds of several prototype GPS sources to be $\\sim 0.1h^{-1}c$. These imply dynamical ages of typically $10^{2-3}$ years.  Additional proof for youth comes from analysis of the overall radio spectra of the somewhat larger CSS sources. Murgia et al (1999) show that their spectra can be fitted with synchrotron aging models, implying ages of typically $10^{3-5}$ years.  The work of these authors shows that GPS/CSO sources are very young and most likely the progenitors of large, extended radio sources. This makes them key objects for radio source evolution studies.   \\subsection*{Tools for radio source evolution studies}  Several authors have used number count statistics and linear size distributions to constrain the luminosity evolution of radio sources (Fanti et al. 1995; Readhead et al. 1996, O'Dea \\& Baum 1997). All these studies find an excess of young objects in relation to the number of old, extended radio sources. This over-abundance of GPS and CSS sources has generally been explained by assuming that a radio source significantly decreases in luminosity over its lifetime. In this way, sources are more likely to contribute to flux density limited samples at young than at old age, causing the apparent excess.  However, in addition to their over-abundance, GPS galaxies are found to be significantly more biased towards high redshift than large extended radio galaxies (Snellen \\& Schilizzi, 2000). This is puzzling since classes of sources representing similar objects at different stages of their evolution are expected to have similar birth functions and redshift distributions. Furthermore, it suggests that the interpretation of their number count statistics, which are averaged over a large redshift range, is not so straightforward. We have postulated a simple evolution scenario which can resolve these puzzles. We argue that the luminosity evolution of a radio-loud AGN during its first $10^5$ years is qualitatively very different from that during the rest of its lifetime. This may be caused by a turnover in the density profile of the interstellar/intergalactic medium at the core-radius of the host galaxy, resulting in an increase in luminosity for young, and a decrease in luminosity for old radio-loud AGN with time. Such a luminosity evolution results in a flatter collective luminosity function for the young objects, causing their bias towards higher redshifts, and their over-abundance at bright flux density levels (Snellen et al. 2000).  An alternative explanation is that GPS sources are indeed young AGN, but mainly short-lived objects, which will never evolve into extended radio sources (Readhead et al. 1994). In that case, the two populations are not directly connected, and no similar cosmological evolution or redshift distribution is necessary. ",
        "conclusions": "GPS galaxies and CSO are now identified as classes of young radio sources. They form a key element in the investigation of the evolution of radio-loud AGN. We report on VSOP observations of 11 and 8 bright GPS sources at 5.0 and 1.6 GHz frequency respectively. First analysis indicates high brightness temperatures consistent with synchrotron self absorption as the cause of their spectral turnover. Comparison with ground-based VLBI datasets taken at earlier epochs confirm the very young ages for the most compact GPS galaxies of a few hundred years."
    },
    "0002/hep-ph0002086_arXiv.txt": {
        "abstract": " ",
        "introduction": "Neutrino physics has become very popular during the past decade. The main impetus in this field has come from the strong improvements in detecting neutrinos; there is now strong evidence from two natural neutrino sources, the solar core \\cite{solarexp} and the Earth's atmosphere \\cite{Atmos}, that neutrinos of a certain flavor disappear on their way to the detector. Furthermore, a neutrino beam produced in the laboratory has been found to change flavor in the LSND experiment \\cite{99LSND}.  The most natural explanation for these anomalies is achieved by extending the Standard Model with neutrino masses. Together with the Cabibbo-Kobayashi-Maskawa (CKM) matrix for the lepton sector, the masses imply flavor mixing; neutrino oscillations have been born \\cite{P57}. Such oscillations can be fully parameterized by the differences between the squared masses $\\dm_{ij}\\equiv m_j^2-m_i^2$ and the mixing angles $\\theta_{ij}$. Also, a possible phase in the CKM matrix could induce CP violation. With these neutrino oscillations, the experiments can be explained by the solutions presented in Table \\ref{noexp}.  However, with this straightforward extension of the Standard Model it is not possible to explain all three experiments simultaneously, simply because the trivial condition \\be \\sum \\dm_{ij} = (m_3^2-m_2^2) + (m_2^2-m_1^2) + (m_1^2-m_3^2) = 0 \\ee cannot be fulfilled by the data. Hence, many interesting new models have been suggested to explain the three different experiments. One of the most appealing possibilities is the existence of a fourth neutrino. Since the $Z^0$ decay width \\cite{PDG98a} excludes more than three standard neutrinos, this new flavor must be inert with respect to the standard weak interaction.  \\begin{table}[ht] \\[\\begin{array}{llll} \\hline \\hline {\\rm Experiment} & {\\rm Favored~Channel} & |\\dm| [\\eV^2] & \\sin^22\\theta \\\\ \\hline {\\rm Solar} &&& \\\\ ~~~{\\rm Vacuum} &\\nu_{\\rm e}\\ra {\\rm anything} & (0.5\\mbox{--}8)\\times 10^{-10} & 0.5\\mbox{--}1 \\\\ ~~~{\\rm MSW~(small~angle)} &\\nu_{\\rm e}\\ra {\\rm anything} & (0.4\\mbox{--}1)\\times 10^{-5} & 10^{-3}\\mbox{--}10^{-2} \\\\ ~~~{\\rm MSW~(large~angle)} &\\nu_{\\rm e}\\ra \\nu_\\mu ~{\\rm or}~ \\nu_\\tau & (3\\mbox{--}30)\\times 10^{-5} & 0.6\\mbox{--}1 \\\\ {\\rm Atmospheric} & \\nu_\\mu \\ra \\nu_\\tau & (1\\mbox{--}8)\\times 10^{-3} & 0.85\\mbox{--}1 \\\\ & \\nu_\\mu \\ra \\nu_{\\rm s} & (2\\mbox{--}7)\\times 10^{-3} & 0.85\\mbox{--}1 \\\\ {\\rm LSND} & \\bar \\nu_\\mu \\ra \\bar \\nu_{\\rm e} & 0.2\\mbox{--}10 & (0.2\\mbox{--}3)\\times 10^{-2} \\\\ \\hline\\hline \\end{array} \\] \\caption{{\\it Results from neutrino oscillation experiments \\cite{raffi99a}. The values are nominal $2\\sigma$ ranges.}} \\label{noexp} \\end{table}  Although the hypothesis of such a sterile neutrino may seem highly speculative, its possible existence is the most far-reaching implication of the current experimental situation. It is certainly worthwhile to investigate its consequences! Sterile neutrinos would have a strong impact on certain astrophysical environments due to their mixing with active neutrinos. These effects would become especially important under extreme conditions, like in core-collapse supernovae \\cite{ed} and in the Early Universe.  It is the impact of sterile neutrinos in the Early Universe, and in particular the outcome of the primordial nucleosynthesis, that has been the subject of an intense recent debate. By incorporating sterile neutrinos into the well understood mechanism of standard Big Bang nucleosynthesis (BBN), two frontiers of physics would be connected in one phenomenon, making this a powerful tool for predictions.  The production of primordial nuclei depends sensitively on the expansion rate of the Universe, which in turn depends on the number density of sterile neutrinos. This simple correlation is well known \\cite{KTTEU}; the actual challenge is to deduce how much the sterile neutrino sector is populated. Since the only connection between sterile neutrinos and ordinary matter, apart from gravity, is given by oscillations, this phenomenon must be the key for creating a relation between sterile neutrino parameters and Big Bang nucleosynthesis predictions.  In the beginning of the 1990's, the number density for sterile neutrinos at the time of primordial nucleosynthesis was calculated depending on the parameters for oscillations between an active and a sterile neutrino, $\\dm$ and $\\theta$ \\cite{EKMBD90f,BD91a}. The idea was to constrain the allowed parameter space, since strong observational results (which have weakened since) set an upper bound on the number density of the $\\nu_{\\rm s}$. In these calculations, the influence of the thermal plasma on the mixing parameters was included, but the small CP asymmetric contributions were neglected. The constraints were so strong that they clearly excluded the $\\nu_\\mu\\leftrightarrow\\nu_{\\rm s}$ solution for the atmospheric neutrino anomaly.  In 1995, Foot, Thomson and Volkas \\cite{FTV96a} found an interesting effect by including the CP asymmetric contributions; the asymmetry could induce different population rates for the sterile neutrinos and antineutrinos, or equivalently, different depopulation rates for the active neutrinos and antineutrinos. This would change the neutrino asymmetry and thus would have an influence on the population rates. They found that this back-reaction could amplify an initially small CP asymmetry by several orders of magnitude for a large range of parameter space. (Actually, this back-reaction had been discussed much earlier by Barbieri and Dolgov \\cite{BD91a}, but they erroneously found the effect to be small.)  The implications from this new effect were manifold. The most interesting conclusion was given by a model in which $\\nu_\\tau\\leftrightarrow\\nu_{\\rm s}$ oscillations created a large $\\nu_\\tau$ asymmetry, which in turn suppressed $\\nu_\\mu\\leftrightarrow\\nu_{\\rm s}$ oscillations [12--17]. % In such a model, the atmospheric neutrino anomaly could be explained by $\\nu_\\mu\\leftrightarrow\\nu_{\\rm s}$ oscillations without the sterile neutrinos coming into thermal equilibrium. Another thoroughly discussed implication treated the impact of a large $\\nu_{\\rm e}$ asymmetry on primordial nucleosynthesis \\cite{FV96a,FV97a,BFV98a}; the large chemical potential $\\mu_{\\rm e}$ would change the neutron-to-proton ratio and thus the nuclear abundances.  Clearly, sterile neutrinos can only be included correctly into BBN if the mechanism producing large neutrino asymmetries is revealed. Unfortunately, this has proven to be a very difficult task. A series of papers have been published on this subject by several groups [12--36], % but the situation remains unclear. The main problem is that the system of differential equations describing the mixing is very complex: one has to treat neutrinos of different momentum separately, and one especially has to take care of the non-linear terms in the equations. Therefore, most works tried to solve the problem with numerical calculations. Thereby, they often used either the adiabatic approximation, which simplifies the equations for a given momentum, or they neglected the momentum distribution. Only during the past year some authors claim to have solved the full system of differential equations numerically \\cite{SF98a,F98a,FV99a}.  Due to these difficulties in solving the system, the different works present contradictory solutions: when neglecting the momentum distribution, the neutrino asymmetry shows an oscillating behavior \\cite{S96a,EKS99a}. The system then ends up with a calculable value of the $\\nu$ asymmetry, but with unpredictable sign. Such a scenario would introduce domains in the Early Universe with different signs of neutrino asymmetry \\cite{SF99a}. On the other hand, applying the adiabatic approximation \\cite{FV96a} is questionable since the neutrino oscillations are not adiabatic close to the resonance for the parameter space of interest. Here the neutrino asymmetry does not oscillate. Since the numerical calculation using the full system of differential equations is very CPU-time consuming, one also has to question whether the calculations are done with sufficient accuracy and whether the numerics are stable at all. After all, non-linear systems tend to show chaotic behavior. Besides, different results achieved by these calculations are contradictory on the point of oscillating neutrino asymmetry \\cite{F98a,SF99a}.  All these numerical works have in common that they predict a similar final absolute neutrino asymmetry (if effects of neutrino domains are ignored) of orders $10^{-2}\\mbox{--}1$. But even this outcome has been questioned very recently in an analytical work by Dolgov {\\it et al.} \\cite{SEMIKOX}, who end up with a final asymmetry several orders of magnitude lower than the former results. Thus, the only point on which all works agree is that oscillations between sterile and active neutrinos in the Early Universe have the potential of creating a neutrino asymmetry at least of order $10^{-5}$, a number which is still large compared to the baryon asymmetry, $\\eta=\\mathcal{O}(10^{-9})$.  The main problem is to get a grip on the exact behavior of the non-linear system. Our work will try to bring some clarity into this subject. Most of the former papers were strongly based on numerical calculations. In this paper, the discussion will be based on a thorough analytical treatment, which will be supported by simple numerical calculations. We will describe the evolution of the system, neglecting the momentum distribution, and compare our results with a numerical solution for the parameters $(\\dm,\\sin 2\\theta_0)=(-1 \\eV^2,10^{-4})$. It should be the job of future investigations to include the effects from the momentum distribution.  In Chapter 2, those aspects of the Early Universe and Big Bang nucleosynthesis will be summarized which are important for our system. Chapter 3 contains all relevant information about neutrino oscillations in matter. In Chapter 4, we analyze the system in the simple two flavor case. After introducing initial conditions, we prove that our system creates oscillations of the lepton asymmetry. In Chapter 5, we summarize our findings.  \\chapter{Big Bang Nucleosynthesis}  We study the influence of neutrinos, in particular their effective number of flavors and the $\\nu_{\\rm e}$ chemical potential, on Big Bang nucleosynthesis (BBN). To this end we first describe the expansion of the Early Universe before and during BBN. We define number density and lepton asymmetry. Next, we give a short summary of BBN, concentrating on the main product, \\Hec, and we estimate the influence the neutrino sector can have on its abundance. Finally, we summarize the observational results. ",
        "conclusions": "We have shown that a deviation of the effective number of neutrinos of order 1, as well as a $\\nu_{\\rm e}$ asymmetry exceeding about $10^{-2}$, will have a measurable effect on the outcome of BBN. Sterile neutrinos have the potential to change both $\\Nnueff$ and $\\Anue$. If we knew how these two variables depend on the mixing parameters of the sterile neutrinos, we could use the measured primordial element abundances to derive constraints on the mixing parameters. On the other hand, if the existence of sterile neutrinos was proven by future experiments such as MiniBooNE \\cite{MINIBOONE}, a detailed understanding of their impact on primordial nucleosynthesis would be necessary to constrain the free parameters of BBN, in particular the baryon asymmetry $\\eta$.  \\chapter{Neutrino Oscillations}  Neutrino oscillations play a crucial role in our considerations of the Early Universe involving sterile neutrinos. In this chapter, we derive the density matrix formalism for neutrino oscillations in media between any two neutrino flavors, active or sterile. This includes the direct influence of matter on the vacuum neutrino oscillations as well as scattering processes which tend to destroy the coherence of the oscillations. Our analysis will only be applicable for temperatures between a few$\\MeV$ and $100\\MeV$ and for neutrino asymmetries $A_\\nu\\ll1$."
    },
    "0002/astro-ph0002116_arXiv.txt": {
        "abstract": "The extremely energetic ($\\rm \\sim 10^{-4} erg/cm^2$) gamma--ray burst (GRB) of 1999 December 8 was triangulated to a $\\sim$ 14 sq. arcmin. error box $\\sim$ 1.8 d after its arrival at Earth with the 3rd interplanetary network (IPN), consisting of the \\it Ulysses, Near Earth Asteroid Rendezvous \\rm (NEAR), and \\it Wind \\rm spacecraft.  Radio observations with the Very Large Array $\\sim$ 2.7 d after the burst revealed a bright fading counterpart whose position is consistent with that of an optical transient source whose redshift is z=0.707. We present the time history, peak flux, fluence, and refined 1.3 sq. arcmin. error box of this event, and discuss its energetics.  This is the first time that a counterpart has been found for a GRB localized only by the IPN. ",
        "introduction": "Many gamma-ray burst counterparts have now been identified using the rapid, precise localizations available from the BeppoSAX spacecraft , as well as from the \\it Rossi X-Ray Timing Explorer \\rm, starting with GRB970228 (Costa et al. 1997; van Paradijs et al. 1997). Such detections occur at a low rate ($\\sim 8 \\, y^{-1}$), and they have been limited to the long--duration events so far, but they have confirmed the cosmological origin of at least this class of bursts. Since 1977, interplanetary networks of omnidirectional GRB detectors have provided precise triangulations of both short and long bursts at rates up to $\\sim$ 1/week, but often the networks have been incomplete, or the data return from the interplanetary spacecraft has been slow.  The present, 3rd IPN is now complete with \\it Ulysses \\rm and NEAR as its distant points (Cline et al. 1999) and, in conjunction with numerous near-Earth spacecraft, can produce precise GRB error boxes within $\\sim$ 1 d, making them useful for multi-wavelength follow-up observations.  Here we present the observations of GRB991208, which was rapidly localized to a small error box, leading to the identification of its radio and optical afterglow, and eventually to the measurement of its redshift. ",
        "conclusions": "For a 25 -- 1000 keV fluence of $\\rm 10^{-4}\\, erg\\, cm^{-2}$ and a redshift z=0.707, the isotropic energy of this burst would have been $\\rm 1.3 \\times 10^{53} erg$; for a peak flux of $\\rm \\sim 10^{-5}\\, erg\\, cm^{-2}\\, s^{-1}$ in the same energy range, the isotropic peak luminosity would have been $\\rm 1.3 \\times 10^{52} erg\\, s^{-1}$.  These estimates assume $\\rm \\Omega=0.2, \\Lambda=0, and \\, H_0=65 \\, km\\, s^{-1}\\, Mpc^{-1}$.  Rhoads (1997) has pointed out that one signature of beaming is a steep decay in the afterglow light curve, $\\propto t^{-2}$.  As the initial optical light curve for GRB991208 indeed appears to decay this steeply, the emission may well be beamed, reducing these estimates.  The current IPN now consists of \\it Ulysses \\rm and NEAR in interplanetary space, and numerous near-Earth spacecraft such as \\it Wind \\rm, \\it BeppoSAX \\rm, and the \\it Compton Gamma-Ray Observatory \\rm.  The \\it Mars Surveyor 2001 Orbiter \\rm will join the network in mid-2001.  The IPN currently observes $\\sim$ 1 -- 2 GRBs per week and is localizing many of them rapidly to small error boxes.  These events tend to be the brighter ones, but apart from this, there is no bias towards any particular event duration; indeed, the IPN generally obtains its smallest error boxes for the short bursts.  Neither is there any sun-angle restriction for the event locations, which means that bursts will be detected whose locations are close to the Sun, as this one was, making prompt radio observations of these positions important.  The other advantages of radio observations over optical are the longer lifetime of the radio afterglow, the immunity from weather, and the freedom to operate at any part of the diurnal cycle. This should increase the rate of counterpart detections substantially over the next several years.   KH acknowledges support for Ulysses operations under JPL Contract 958056, for IPN operations under NASA LTSA grant NAG5-3500, and for NEAR operations under the NEAR Participating Scientist program.  On the Russian side, this work was partially supported by RFBR grant \\# 99-02-17031. We are grateful to R. Gold and R. McNutt for their assistance with the NEAR spacecraft.  RS is supported by NASA grant NCC5-380.  We are indebted to T. Sheets for her excellent work on NEAR data reduction.  Special thanks also go to the NEAR project office for its support of post-launch changes to XGRS software that made these measurements possible.  In particular, we are grateful to John R. Hayes and Susan E. Schneider for writing the GRB software for the XGRS instrument and to Stanley B. Cooper and David S. Tillman for making it possible to get accurate universal time for the NEAR GRB detections."
    },
    "0002/astro-ph0002320_arXiv.txt": {
        "abstract": "A distinct sub-class of anomalous velocity \\hi emission features has emerged from recent high quality surveys of the Local Group environment, namely the compact high velocity clouds (CHVCs). A program of high-resolution imaging with the Westerbork array and the Arecibo telescope has begun to provide many insights into the nature of these objects. Elongated core components with a velocity gradient consistent with rotation (V$_{Rot}\\sim$ 15 \\kms) are seen in many objects. Comparison of volume and column densities has allowed the first distance estimates to be made (600$\\pm$300~kpc). The objects appear to be strongly dark-matter dominated with dark-to-gas mass ratios of 30--50 implied if the typical distance is 700~kpc. ",
        "introduction": "The nature of anomalous velocity \\hi discovered during $\\lambda$21~cm surveys of the Galaxy has been the subject of much debate in the past three decades. There appear to be at least four reasonably distinct classes of this gas, namely: (1) Localized outflows and subsequent inflows associated with massive star formation which are commonly referred to as a ``galactic fountain'' (Shapiro \\& Field, 1976; Bregman, 1980). (2) The stream of tidal debris from the interaction of the Galaxy with the Magellanic Clouds is another major source of anomalous velocity \\hi\\/. Recent imaging (Putman \\& Gibson, 1999) has begun to reveal the true extent of this debris system, and allow it to be distinguished from other components. (3) The extended HVC complexes (named A, C, H, M, \\dots) which span ten's of degrees on the sky and have recently been determined to have both nearby distances of about 10~kpc (Van Woerden et al. 1999) and rather low ($\\sim$0.1 solar) metal abundance (Wakker et al. 1999). These diffuse \\hi structures appear to be currently merging with the Galaxy and each account for some 10$^7$ M$\\odot$ of fresh gas. And finally (4) the system of compact high velocity clouds (CHVCs) cataloged by Braun \\& Burton (1999).  The CHVCs form a distinct class of compact, high-contrast \\hi emission features at anomalous velocity, which can be readily distinguished from the more diffuse components in the recent high quality surveys; the Leiden/Dwingeloo Survey in the North (Hartmann \\& Burton 1997) and the Parkes Multibeam Survey in the South (see Putman \\& Gibson 1999). Their average angular size is only 1~degree FWHM and the total velocity width 30~\\kms~ FWHM. The 65 confirmed members of the CHVC class share a well-defined kinematic pattern. A global search for the reference frame that minimizes the line-of-sight velocity dispersion of the system, returns the Local Group Standard of Rest with high confidence. In this system the velocity dispersion of the population is only 70~\\kms, although it is in-falling toward the Local Group barycenter at about 100~\\kms. These properties suggest that the objects are: (1) associated with the Local Group rather than the Galaxy as such, (2) that they are likely to reside at quite substantial distances and (3) have as yet had little tidal interaction with the more massive Local Group members.  \\begin{figure} \\plotone{braunr1.eps}  \\caption{{\\bf A:}~ Westerbork image of CHVC\\,$204\\!+\\!30\\!+\\!075$ showing \\NH at 1~arcmin resolution; contours are drawn at levels of 20, 50, 100, 200, and $300 \\times 10^{18}$~cm$^{-2}$.  {\\bf B:}~ Intensity--weighted line--of--sight velocity, with contours of $v_{\\rm LSR}$ drawn in steps of 5 \\kms~ from 40 to 85 \\kms.  {\\bf C:} Rotation velocities fit to the two principal components of CHVC\\,$204\\!+\\!30\\!+\\!075$.  The solid lines show the rotation curves of Navarro, Frenk, \\& White (1997) cold--dark--matter halos of the indicated mass and 0.7~Mpc distance.  {\\bf D: top}~ Equilibrium temperature curves for \\hi in an intergalactic environment characterized by a metallicity of 10\\% of the solar value and a dust--to--gas ratio of 10\\% of that in the solar neighborhood, calculated for two values of the neutral shielding column depth, namely $10^{19}$~cm$^ {-2}$ (solid line) and $10^{20}$~cm$^{-2}$ (dashed line).  {\\bf D: bottom}~ Mass volume density, in units of hydrogen nuclei per cubic centimeter, as a function of radius for an NFW cold dark matter halo.  } \\end{figure} ",
        "conclusions": ""
    },
    "0002/astro-ph0002050_arXiv.txt": {
        "abstract": "Spergel \\& Steinhardt proposed the possibility that the dark matter particles are self-interacting, as a solution to two discrepancies between the predictions of cold dark matter models and the observations: first, the observed dark matter distribution in some dwarf galaxies has large, constant-density cores, as opposed to the predicted central cusps; and second, small satellites of normal galaxies are much less abundant than predicted. The dark matter self-interaction would produce isothermal cores in halos and expel the dark matter particles from dwarfs orbiting in large halos. Another consequence of the model is that halos should become spherical once most particles have interacted. Several observations show that the mass distribution in relaxed clusters of galaxies is elliptical. Here, I discuss in particular gravitational lensing in the cluster MS2137-23, where the ellipticity of the dark matter distribution can be measured to a small radius, $r\\sim 70$ kpc, suggesting that most dark matter particles in clusters outside this radius do not collide during the characteristic age of clusters. If true, this implies that any dark matter self-interaction with a cross section independent of velocity is too weak to have affected the observed density profiles in the dark-matter dominated dwarf galaxies, or to have facilitated the destruction of dwarf satellites in galactic halos. If $s_x$ is the cross section and $m_x$ the mass of the dark matter particle, then $s_x/m_x < 10^{-25.5} \\cm^2/\\gev$. ",
        "introduction": "The Cold Dark Matter (CDM) model of structure formation in the universe has been tremendously successful in accounting for a huge variety of available observations (e.g., the Cosmic Background fluctuations, the abundances of clusters of galaxies, peculiar velocity fields, the Ly$\\alpha$ forest), provided that the mean density of matter is only a fraction $ \\Omega_m \\simeq 0.3$ of the critical density, and the existence of vacuum energy with a negative pressure equation of state is allowed to make the universe spatially flat (e.g., \\cite{kp00}; \\cite{ptw99}; \\cite{bops99}; \\cite{sw95}; \\cite{ecf96}; \\cite{crofta99}).  A possible problem of this model has emerged when comparing the density profiles of dark matter halos predicted in numerical simulations, with observations of the rotation curves in dwarf galaxies (\\cite{m94}; \\cite{fp94}; \\cite{nfw96}; \\cite{moorea98}; \\cite{kkbp98}; \\cite{moorea99}). Whereas the observations show linearly rising rotation curves out to core radii greater than $1 \\kpc$ in certain dwarf galaxies where the density is dominated by dark matter everywhere (indicating that the dark matter has a constant density core), the simulations predict that the collapse of collisionless particles of cold dark matter produces cuspy halo density profiles, with a logarithmic slope $-d\\log\\rho / d\\log r > 1$ down to the smallest resolved radius. A second problem is that the number of dwarf galaxies observed in the Local Group is much smaller than the total number predicted from numerical simulations (\\cite{kkvp99}; \\cite{moorea99}).  A solution to this discrepancy has been proposed by Spergel \\& Steinhardt (2000): if the dark matter is self-interacting, with large enough cross section to make most particles in the inner core of a dwarf galaxy interact among themselves over a Hubble time, then an isothermal core will be produced. A clear prediction of this hypothesis is that when most of the particles of a halo within some radius $r_c$ have interacted, then the halo should be close to spherical inside $r_c$, or else be supported by rotation, because the velocity dispersion tensor should become isotropic. This paper examines the consequence of this prediction for the inner parts of rich clusters of galaxies, where highly magnified images of background galaxies are occasionally observed. We will find that severe restrictions on the collisional dark matter hypothesis are obtained. ",
        "conclusions": "The modeling of multiple images of background galaxies produced by gravitational lensing in clusters of galaxies require elliptical models of the mass distribution in order to reproduce their positions and magnifications successfully (\\cite{kmfm93}; \\cite{mfk93}; \\cite{kneiba95}). The last section discussed the specific example of MS2137-23, where the misalignment in the position of two images relative to the cluster center can be used to constrain the ellipticity in a model-independent way: the ellipticity of the dark matter halo around the central galaxy must be greater than $0.1$ within the image $I$, which is at $22\\arcsec 5$ from the cluster center, corresponding to a distance of $65 h^{-1} \\kpc$. The fact that the dark matter halos of galaxy clusters are elliptical within this small radius implies that the dark matter particles have not collided over the age of the cluster. As shown in \\S 2, this also implies that the observed cores of the dark matter halos in dwarf galaxies are too big to have been caused by dark matter self-interaction, as proposed by Spergel \\& Steinhardt (2000).  Further evidence supporting that cluster dark matter halos are elliptical at radii $\\sim 100$ kpc comes from the similarity with the ellipticity of the optical isophotes of the central cluster galaxies in both the magnitude of the ellipticity and the orientation of the major axis (\\cite{mfk93}; \\cite{kmfm93}; \\cite{kneiba95}). If the underlying dark matter distribution became spherical due to the collisions, the ellipticity of the stellar distribution would be reduced (although not eliminated, owing to the anisotropy in the velocity dispersion tensor). According to Hammer \\etal (1997), the central galaxy in MS2137-23 has ellipticity $\\epsilon = 0.16 \\pm 0.02$ beyond the radius of the radial arc, and the best fit ellipticity for the lens model is $\\epsilon = 0.18$ (see also \\cite{kneiba95} for similar conclusions obtained in the cluster A2218). We note again that the ellipticities of the optical isophotes decline at a radius smaller than that probed by gravitational lensing (Porter \\etal 1991).  The ellipticity of the cluster halo can be used to place an upper limit on the interaction rate of the dark matter, in terms of the cross section $s_x$ and mass $m_x$ of the dark matter particle. We assume here that the collisional radius must be smaller than the distance from the center to the long tangential arc and two other arclets (these images are A01-A02, A2 and A4 in \\cite{hammera97}, and they also require an ellipticity similar to that of the central galaxy in the lensing models), which is about $70 \\kpc$. The dark matter density at this radius is $\\rho \\simeq \\Sigma_{crit}/2r$, where the critical surface density is $\\Sigma_{crit}\\simeq 1\\gm\\cm^{-2}$ for a source at $z_s=1$. Assuming also a cluster velocity dispersion $\\sigma=1000 \\kms$ (roughly the minimum value required given the Einstein radius of the cluster), and a cluster age $t_c = 5\\times 10^9$ years, we obtain the upper limit \\begin{equation} {s_x \\over m_x} < {1 \\over \\rho\\, 2^{1/2}\\sigma\\, t_c } \\simeq 10^{-25.5} { \\cm^2 \\over m_p } \\simeq  0.02 {\\cm^2 \\over \\gm} ~. \\end{equation} For the dwarf galaxies DDO 154, DDO 170, and DDO 236 mentioned in \\S 2, with velocity dispersion $\\sigma=(28,52,45) \\kms$, and core radii $(3,2.5,6) \\kpc$, the time it would take for the collisional radius to reach the value of their core radii if $s_x/m_x$ were equal to the above upper limit is $t = (40,5,40)\\times 10^{10}$ years, respectively [where we have used the relation $t \\propto \\sigma^3/r_c^2$, from eq.\\ (\\ref{colprop}) ].  The limit we have obtained on the self-interaction of the dark matter also rules it out as an explanation for the low abundance of dwarf galaxies in the Local Group, compared to the predictions of halo satellites abundances from numerical simulations (\\cite{kkvp99}; \\cite{mooreb99}). In order to strike out the dark matter particles, the satellite halos must be moving in an orbit inside the collisional radius. For example, in the Milky Way halo (with $\\sigma \\simeq 150 \\kms$), the collisional radius cannot be greater than about $6 \\kpc$, if $r_c < 100 \\kpc$ in a cluster with $\\sigma = 1000 \\kms$ (where we use the scaling $r_c \\propto \\sigma^{3/2}$).  Finally, we mention three ways by which the collisional dark matter hypothesis might still remain viable as an explanation of the constant density cores observed in some dwarf galaxies. A first possibility is that the presence of substructure in the mass distribution of MS2137-23, or of other massive structures projected on the line of sight of the cluster, introduces an external shear that would modify the positions of the images. However, this seems unlikely as discussed in \\S 3, because elliptical models fit the observed positions and shapes of the images remarkably well with fewer model parameters than observational constraints, and an external shear induces a lensing potential different than a constant ellipticity, and would not generally be aligned with the major axis of the galaxy. The second possibility is that the ellipticity of the dark matter could be supported by rotation, instead of anisotropic velocity dispersion. However, halos formed by collisionless collapse are known to rotate very slowly (\\cite{be87}; \\cite{wqsz92}), and the collisions would further slow down the rotation of the central parts of the halo by enforcing solid body rotation. Finally, there is the possibility that the cross section for the dark matter interaction decreases with velocity. Here we have assumed the cross section to be constant; if it were proportional to $v^{-1}$ (see, e.g., Firmani \\etal 2000), then the constraints we have used here from gravitational lensing in clusters of galaxies would allow a large enough collisional radius in dwarfs to explain their dark matter core radii."
    },
    "0002/astro-ph0002099_arXiv.txt": {
        "abstract": "A change of the component separation in the profiles of the binary pulsar PSR B1913+16 has been observed for the first time (Kramer 1998) as expected by geodetic precession. In this work we extend the previous work by accounting for recent data from the Effelsberg 100-m telescope and Arecibo Observatory and testing model predictions. We demonstrate how the new information will provide additional information on the solutions of the system geometry. ",
        "introduction": "The binary pulsar PSR B1913+16 exhibits a measurable amount of geodetic precession due to a misalignment of its spin axis and orbital angular momentum (relativistic spin orbit coupling) by an angle $\\delta$ (see also Stairs et al., for geodetic precession in PSR B1534+12). The angle between the line of sight and the spin axis, $\\zeta$, changes according to: \\begin{displaymath} \\cos\\zeta(t) = -\\cos\\delta\\cos i +\\sin\\delta\\sin i \\cos \\Omega_{{\\rm prec}} (t-T_0) \\;, \\end{displaymath} where $i$ is the orbital inclination of the binary system. General Relativity (GR) predicts a rate of precession $\\Omega_{{\\rm prec}}=1.21$ deg yr$^{-1}$, leading to a change in the pulse profile and polarization characteristics with time. \\\\ Previous observations by Weisberg, Romani and Taylor (1989) detected a change in the relative amplitude of the trailing and leading component, but no change in the component separation. Also, there were no changes in the position angle swing of the linear polarization noted (Cordes, Wasserman \\& Blaskiewicz 1990). ",
        "conclusions": ""
    },
    "0002/astro-ph0002266_arXiv.txt": {
        "abstract": "In this paper, we present a multi-color photometric study of the nearby spiral galaxy M81, using images obtained with the Beijing Astronomical Observatory $60/90$ cm Schmidt Telescope in 13 intermediate-band filters from 3800 to 10000{\\AA}.  The observations cover the whole area of M81 with a total integration of 51 hours from February 1995 to February 1997. This provides a multi-color map of M81 in pixels of $1\\arcsec.7 \\times 1\\arcsec.7$.  Using theoretical stellar population synthesis models, we demonstrate that some BATC colors and color indices can be used to disentangle the age and metallicity effect.  We compare in detail the observed properties of M81 with the predictions from population synthesis models and quantify the relative chemical abundance, age and reddening distributions for different components of M81.  We find that the metallicity of M81 is about $Z=0.03$ with no significant difference over the whole galaxy.  In contrast, an age gradient is found between stellar populations of the central regions and of the bulge and disk regions of M81: the stellar population in its central regions is older than 8 Gyr while the disk stars are considerably younger, $\\sim 2$ Gyr.  We also give the reddening distribution in M81. Some dust lanes are found in the galaxy bulge region and the reddening in the outer disk is higher than that in the central regions. ",
        "introduction": "Spatially resolved information about the age, metallicity and interstellar medium reddening of galaxies is a powerful tool to study galaxy evolution since it provides essential clues in star formation history, chemical composition and enrichment history and environment of galaxies (\\cite{Buzzoni89}).  To obtain such information, we need to know the stellar component and the overall properties of stellar populations (\\cite{Leitherer95}). Ideally one would like to study resolved individual stars in galaxies. However, given the limited spatial resolution of current telescopes, this is only possible for a few very nearby galaxies. As a result, the stellar content of even some relatively simple galaxies remain to be unraveled (\\cite{Thuan91}). Information of stellar population and star formation history in these galaxies, however, can still be obtained from studying the {\\it integrated} properties of the stars (e.g. \\cite{Schmitt96}, \\cite{Goerdt98}).  Since the pioneering work of Tinsley (1972) and Searle et al. (1973), evolutionary population synthesis has become a standard technique to study the stellar populations of galaxies. This is a result of improvements in the theory of the chemical evolution of galaxies, star formation, stellar evolution and atmospheres, and the development of synthesis algorithms and the availability of various evolutionary synthesis models. A comprehensive compilation of such models was published by Leitherer et al. (1996) and Kennicutt (1998). Widely used models include those from the Padova and Geneva group (e.g. \\cite{Schaerer97}, \\cite{Schaerer98}, \\cite{Bressan96}, and \\cite{Chiosi98}), GISSEL96 (\\cite{Charlot91}, \\cite{Bruzual93}, \\cite{Bruzual96}), PEGASE (\\cite{Fioc97}) and STARBURST99 (\\cite{Leitherer99}).  Many previous studies of integrated stellar populations use spectroscopic data, usually for limited regions in galaxies.  In this paper, we will, instead, use multi-color photometry to probe the stellar populations.  The multi-color photometry provides accurate spectral energy distributions (SEDs) for the {\\it whole} galaxy, although at low spectral resolution. We shall demonstrate that it is a powerful tool to study the structure and evolution of the galaxy together with the theoretical evolutionary population synthesis methods (for an application of a similar technique with, but with fewer colors, to moderate redshifts, see Abraham et al. 1999). For this purpose, we pick M81 as the first application of this multi-color approach.  M81 is an excellent candidate because it is a nearby early-type Sab spiral galaxy at a distance of 3.6 Mpc and with an angular size of $\\sim 26^{\\prime}$. The angular extent is large enough such that the disk and bulge regions are well separated from the ground. It has been the subject of numerous previous studies providing a wealth of information with which to compare the new metallicity and internal reddening distribution. The internal reddening has been studied by Kaufman et al. (1987, 1989), Devereux et al. (1995), Ho et al. (1996) and Allen et al. (1997). The metallicity has been studied by Stauffer \\& Bothun (1984), Garnett \\& Shields (1987), Brodie \\& Huchra (1991) and Perelmuter et al. (1995). In this paper, we present a further detailed study of M81 using the unique dataset obtained from the BATC \\footnote[1]{The Beijing-Arizona-Taiwan-Connecticut Multicolor Sky Survey} multi-color sky survey.  The outline of the paper is as follows.  Details of observations and data reduction are given in section 2. In section 3, we provide a brief description of the  model, and analyze the evolution of the integrated colors, color indices with age and metallicity. The observed two-dimensional spectral energy distributions (SEDs) of M81 were analyzed using stellar population synthesis models of Bruzual \\& Charlot (1996). The distributions of metallicity, age and interstellar reddening are given in section 4. In section 5, we discuss how different star formation histories and stellar population synthesis models change our results, and compare our results with previous studies. . The conclusions are summarized in section 6. ",
        "conclusions": "In this paper, we have, for the first time, obtained a two-dimensional SED of M81 in 13 intermediate colors with the BAO 60/90 cm Schmidt telescope. Below, we summarize our main conclusions.  \\begin{itemize}  \\item Using the new extensive grid of GSSPs covering a wide range of metallicity and age, we calculated the colors and color indices for 13 colors in BATC intermediate-band filter system. We find that some of them can be used to break the age and metallicity degeneracy, which enables us to obtain two-dimensional maps of metallicity, interstellar reddening and age of M81.  \\item From the two dimensional metallicity distribution of M81, we find no obvious metallicity gradient from the central regions to the outer disk. In most part of M81, the mean metallicity is about 0.03 with variation  $\\lesssim 0.005$. Some regions in M81, however, have higher metallicity; they are mostly located in the spiral arms and around HII regions, where the younger component resides.  \\item From the two dimensional age distribution in M81, we find that the mean ages of the stellar populations in the central regions are older than those in the outer regions, which suggests that star formations in the central regions occurred earlier than the outer regions.   \\item We find a strong difference in reddening between the bulge region and the disk region. In the bulge area, the reddening, E(B-V), is in the range of 0.08 to 0.15.  The mean reddening in the disk area is higher, about 0.2. There are some high reddening spiral-like cirrus in the bulge.  \\item In order to understand how sensitive our method is to different assumptions about star formation history and different stellar population synthesis models, we have studied an exponential star formation history and compared the results obtained with GSSP and PSSP.  We find that although the precise values of age, metallicity and interstellar reddening are different, the general trend of in the metallicity, age and reddening distributions is similar.  \\item  Finally, we have compared the internal reddening and metallicity maps of M81 with previous studies. We find that the agreements are generally good. In addition, we find that the properties for the bright HII regions and other parts may be different.  \\end{itemize}  The results of M81 presented here illustrate that our method and observational data provide an efficient way to study the distribution of metallicity, age and interstellar reddening for nearby face-on galaxies. Similar data have already been collected for similar galaxies M13, NGC589 and NGC5055. The analysis results of these galaxies will be published in a forthcoming paper."
    },
    "0002/astro-ph0002272_arXiv.txt": {
        "abstract": "Using the {\\em RossiXTE} experiment, we detect weak X-ray emission from the recurrent Be/X-ray transient A\\,0535+26 at a time when the optical counterpart V725 Tau displayed H$\\alpha$ in absorption, indicating the absence of a circumstellar disc. The X-ray radiation is strongly modulated at the 103.5-s pulse period of the neutron star, confirming that it originates from A\\,0535+26. The source is weaker than in previous quiescence detections by two orders of magnitude and should be in the centrifugal inhibition regime. We show that the X-ray luminosity cannot be due to accretion on to the magnetosphere of the neutron star. Therefore this detection represents a new state of the accreting pulsar. We speculate that the X-ray emission can be due to some matter leaking through the magnetospheric barrier or thermal radiation from the neutron star surface due to crustal heating. The observed luminosity is probably compatible with recent predictions of thermal radiation from X-ray transients in quiescence. The detection of the X-ray source in the inhibition regime implies a reduced density in the outflow from the Be companion during its disc-less phase. ",
        "introduction": "Be/X-ray binaries are X-ray sources composed of a Be star and a neutron star. Most of these systems are transient X-ray pulsars displaying strong outbursts in which their X-ray luminosity increases by a factor $\\ga 100$ (see Negueruela 1998). In addition, those systems in which the neutron star does not rotate fast enough for centrifugal inhibition of accretion to be effective (see Stella et al. 1986) display persistent X-ray luminosity at a level $\\la 10^{35}$ erg s$^{-1}$. The high-energy radiation is believed to arise due to accretion of material associated with the Be star by the compact object. It has long been known that accretion from the fast polar wind that is detected in the UV resonance lines of the Be primaries cannot provide the observed luminosities, even for detections at the weakest level (see Waters et al. 1988 and references therein; see also the calculations made for X Persei by Telting et al. 1998). Therefore it is believed that the material accreted comes from the dense equatorial disc that surrounds the Be star. Waters et al. (1989) modelled the radial outflow as a relatively slow ($\\sim 100$ km s$^{-1}$) dense wind. However most modern models for Be stars consider much slower outflows, due to strong evidence for rotationally dominated motion (quasi-Keplerian discs). This is due not only to the line shapes (see Hanuschik et al. 1996), which set an upper limit on the bulk motion at $v \\la 3\\:{\\rm km}\\,{\\rm s}^{-1}$ (Hanuschik 2000), but also to the success of the Global One-Armed Oscillation model (which can only work in quasi-Keplerian discs) at explaining V/R variability in Be stars (Hummel \\& Hanuschik 1997; Okazaki 1997ab). The viscous decretion disc model (Okazaki 1997b; Porter 1999; Negueruela \\& Okazaki 2000) considers material in  quasi-Keplerian orbit with an initially very subsonic outflow velocity that is gradually accelerated by gas pressure and becomes supersonic at distances $\\sim 100 R_{*}$, i.e., much further than the orbits of neutron stars in Be/X-ray transients.   The transient A\\,0535+26 is one of the best studied Be/X-ray binaries (Clark et al. 1998 and references therein). It contains a slowly rotating ($P_{{\\rm s}} = 103 \\: {\\rm s}$) neutron star in a relatively wide ($P_{{\\rm orb}} = 110.3 \\: {\\rm d}$) and eccentric ($e = 0.47$) orbit around the B0IIIe star V725 Tau (see Finger et al. 1996; Steele et al. 1998). After its last giant outburst in February 1994 (Clark et al. 1998; Negueruela et al. 1998), the strength of the emission lines in the optical spectrum of V725 Tau has declined steadily. The last normal (periodic) outburst took place in September 1994 and the source has since not been detected by the BATSE experiment on board {\\em CGRO}. ",
        "conclusions": "A\\,0535+26 was active at a time when the optical counterpart V725 Tau showed no evidence for a circumstellar disc. The luminosity was two orders of magnitude lower than in previous quiescence detections and accretion was centrifugally inhibited. Given that the observed luminosity cannot be due to accretion onto the magnetosphere, we are observing either some material leaking through the magnetosphere or thermal emission from the heated core of the neutron star. In any case, this detection represents a state of an accreting X-ray pulsar that had not been observed before.  Further observations of Be/X-ray binaries in a similar state (when their companions have lost their discs and very little material can reach the vicinity of the neutron star) are needed. Observations with {\\em Chandra} or {\\em XMM}, which combine much higher sensitivities with more adequate energy ranges, could determine whether the observed spectrum is compatible with thermal emission models."
    },
    "0002/astro-ph0002044_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "The CMB is a mature subject. The spectral distortions are well understood, and the Sunyaev-Zeldovich effect provides a unique tool for studying galaxy clusters at high redshift. Global distortions will eventually be found, most likely first at very large $l$  due to the cumulative contributions from hot gas heated by radio galaxies, AGN, and galaxy groups and clusters. For gas at $\\sim 10^6 - 10^7$ K, appropriate to gas in galaxy potential wells, the thermal and kinematic contributions are likely to be comparable.  CMB anisotropies are a rapidly developing field, since the 1992 discovery with the COBE DMR of large angular scale temperature fluctuations. At the time of writing, the first acoustic peak is being mapped with unprecedented precision that will enable definitive estimates to be made of the curvature parameter. More information will come with all-sky surveys to higher resolution (MAP in 2000, PLANCK in 2007) that will enable most of the cosmological parameters to be derived to better than a few percent precision if the adiabatic CDM paradigm proves correct. Degeneracies remain in CMB parameter extraction, specifically between $\\Omega_0$, $\\Omega_b$ and $\\Omega_{\\Lambda}$, but these can be removed via large-scale structure observations, which effectively constrain $\\Omega_{\\Lambda}$ via weak lensing. The goal of studying reionization will be met by the interferometric surveys at very high resolution ($l\\sim 10^3 - 10^4$).  Polarization presents the ultimate challenge, because the foregrounds are poorly known. Experiments are underway to measure polarization at the 10 percent level, expected on degree scales in the most optimistic models. However one has to measure polarisation at the 1 percent level to definitively study the ionization history and early tensor mode generation in the universe, and this may only be possible with long duration balloon or space experiments.  CMB anisotropies are a powerful probe of the early universe. Not only can one hope to extract the cosmological parameters, but one should be able to measure the primordial power spectrum of density fluctuations laid down at the epoch of inflation, to within the uncertainties imposed by cosmic variance. In combination with new generations of deep wide field galaxy surveys, it should be possible to unambiguously measure the shape of the predicted peak in the power spectrum, and thereby establish unique constraints on the origin of the large-scale structure of the universe.   \\small"
    },
    "0002/astro-ph0002334_arXiv.txt": {
        "abstract": "\\vskip .2 truein We give distance--redshift relations in terms of elliptic integrals for three different mass distributions of the Friedmann-Lema\\^\\i tre-Robertson-Walker (FLRW) cosmology. These models are dynamically pressure free FLRW on large scales but, due to mass inhomogeneities, differ in their optical properties. They are the filled-beam model (standard FLRW), the empty-beam model (no mass density exists in the observing beams) and the 2/3 filled-beam model. For special $\\OM$--\\ $\\OL$ values the elliptic integrals reduce to more familiar functions. These new expressions for distance-redshift significantly reduce computer evaluation times. ",
        "introduction": "\\label{sec-intro} As limits on the global cosmological parameters $\\OM$ and $\\Lambda$ have been refined, \\cite{SB} and \\cite{PS1}, the optical inadequacy of the standard distance-redshift relation ($D$-$z$) of FLRW has become more apparent. The problem was first recognized long ago by \\cite{Zel}, \\cite{BB}, and \\cite{KR} but the lack of relevant data limited its significance. Even though the average mass density parameter $\\OM$ (along with $H_0$ and $\\Lambda$) determines the large scale dynamic behavior of the pressure free universe, knowledge of the actual mass inhomogeneity is necessary to accurately determine these parameters from most observations. Most observations determine $\\OM$ and $\\Lambda$ by (indirectly) comparing theoretical $D$-$z$ curves to observed data. However, $D$-$z$ depends on more than the average mass density. It can depend significantly on details of how the mass is distributed, \\ie on how inhomogeneous the mass is on the scale of the widths of the observing beams. If some significant fraction ($\\rho_I/\\rho_0 \\le 1$) of the total mass density is in the form of inhomogeneities and is excluded from the lines of sight to the distant objects observed, a modified, \\ie a partially filled-beam $D$-$z$ is  required.  The necessity of taking into account the effect of inhomogeneities on observations is relatively easy to understand. Homogeneous matter inside an observing beam of light gravitationally focuses the beam much differently than does an equal-mass clump of externally lensing matter. The simplest correction for this gravity-light effect requires the introduction of another parameter $\\nu,\\ 0\\le\\nu\\le 2,$ which gives the fraction $\\rho_I/\\rho_0 =\\nu(\\nu+1)/6$ of the mass density of the universe removed from the observing beams as inhomogeneities. Using $\\nu$ rather than  $\\rho_I/\\rho_0$ or some other parameter is dictated by the mathematics of special functions. A reduced mass density in an observing beam causes it to diverge relative to a standard FLRW beam. For an observed object in such a universe to have the standard FLRW angular size it would thus have to be moved to a smaller $z$; i.e., objects will appear less bright than in the standard FLRW universe. A reasonable application of this model to SNe Ia observations takes $\\rho_I$ as the galactic contribution to the total mass density $\\rho_0$ and the remaining contribution as a smooth intergalactic medium. Galaxies are easily excluded from SNe Ia foregrounds by selection (intended or not) and if galaxy mass roughly follows light, including their mass in $\\rho_I$ is appropriate. In the partially filled-beam model where the additional  parameter $\\nu\\ne 0$ has been introduced, only lensing by mass clumps external to the beam has been neglected. To compare individual observations to  $D$-$z$ of this model  requires only  an occasional lensing correction; however, comparison with the standard FLRW $D$-$z$ ($\\nu=0$) model requires a defocusing correction for the partially empty-beam of every observation, as well as the occasional lensing  correction. If only weak and transparent lensing occurs (to the $z_{max}$ being observed) the standard  FLRW $D$-$z$ ($\\nu=0$) should give the mean $D$-$z$ curve. \\cite{WY} argues that by using flux-averaging the mean can be accurately obtained. \\cite{KRa} and \\cite{KRb} claims that determining cosmological parameters from data compared with the partially filled Hubble curves given here is likely to be easier. Beyond selection effects, unknown lensing probabilities can be highly non-Gaussian and should make the mean more difficult to observationally determine, \\ie should require more data if a given accuracy of the cosmic parameters is to be obtained, \\cite{BB,HW,HD}. The down side for partially filled-beam models is that you must select against lensing and must determine the additional parameter $\\nu$.  In Sec.\\,\\ref{sec-lumdist} we outline the procedure required to obtain $D$-$z$ for partially filled-beam FLRW observations and how the result simplifies for the three special cases of $\\nu =$ 0, 1, and 2. In Sec.\\,\\ref{sec-results} we give the new results for these three special cases. Some concluding remarks are given in Sec.\\,\\ref{sec-conclusions} and in the Appendix we discuss our Fortran implementation of these results. ",
        "conclusions": "\\label{sec-conclusions}  We have given useful forms for the luminosity distance in three currently relevant cosmologies. They are all dynamically FLRW cosmologies in the large but differ in how gravitating matter effects optical observations. The models are labeled by an additional parameter $\\nu$ ($\\nu$ = 0, 1, and 2) beyond the familiar $H_0,\\Omega_m,$ and $\\Lambda$. The $\\nu=0$ model is standard FLRW where all matter is homogeneous and transparent on the scale of the observing beam widths. This model is called the `filled-beam' model. The $\\nu=2$ model  assumes the opposite; all matter is inhomogeneous and excluded from the observing beams. This extreme case is called the `empty-beam' model. The $\\nu=1$ model assumes that 1/3 of the mass density of the universe is excluded from observing beams and hence it is the `two-thirds filled-beam' model. These three cases were singled out because their distance-redshift relations can be given in terms of incomplete elliptic integrals; functions which are universally available in computer libraries and very efficiently evaluated.\\footnote{ The results appearing in Section \\ref{sec-results} have been coded and are posted at http://www.nhn.ou.edu$\\sim$thomas/z2dl.html. This code is discussed in the Appendix and compared to the numerical integration times of \\cite{KHS}.} For the $\\nu=1$ and 2 cases, somewhat simpler expressions than what we have given exist, but only for complex arguments of the elliptic integrals. We chose to give expressions whose arguments are real and which can be rapidly evaluated. Results are available for all $0\\le\\nu\\le 2$ but only in terms of the less familiar and unavailable Heun functions, \\cite{KRa}. We have extended the flat space, $\\OO=1$, results given here  to arbitrary filling parameter $\\nu$. These new results will be available shortly. Related results have been independantly found by \\cite{DM}. A calculation similar to the $\\nu=1$ case given here is that of the age of the Universe as a function of redshift and can be found in \\cite{TK}."
    },
    "0002/astro-ph0002102_arXiv.txt": {
        "abstract": " ",
        "introduction": "The Astrophysics Data System (ADS) provides access to the astronomical literature through the World Wide Web (WWW).  It is widely used in the astronomical community.  It is accessible to anybody world-wide through a forms based WWW interface.  A detailed description of the history of the ADS is presented in the ADS Overview article (\\cite{adsoverview}, hereafter OVERVIEW).  The system contains information from many sources (journals, other data centers, individuals).  A detailed description of the data that we get and how they are included in the ADS is presented in the ADS Data article (\\cite{adsdata}, hereafter DATA).  The incoming data are processed and indexed with custom-built software to take advantage of specialized knowledge of the data and the astronomical context.  A description of this processing is given in the ADS Architecture article (\\cite{adsarchitecture}, hereafter ARCHITECTURE).  This article describes the development and the current status of the ADS Abstract Service user interface and search engine.  The ADS was created as a system to provide access to astronomical data (\\cite{1992ald2.proc..387M}).  In 1993 the ADS started to provide access to a set of abstracts obtained from the NASA/STI (National Aeronautics and Space Administration/Scientific and Technical Information) project (\\cite{1993adass...2..132K}).  The user interface was built with the proprietary software system that the ADS used at that time.  The search engine of this first implementation used a commercial database system.  A description of the system at that time is in \\cite{1994ExA.....5..205E}.  In 1994, the World Wide Web (WWW, \\cite{www}) became widely useful through the NCSA Mosaic Web Browser (\\cite{mosaic}).  The design of the ADS Abstract Service with a clean separation between the user interface and the search engine made it very easy to move the user interface from the proprietary ADS system to the WWW.  In February 1994, a WWW interface to the ADS Abstract Service was made available publicly.  The WWW interface to the ADS is described by \\cite{1995adass...4...28E} and \\cite{1995VA.....39..217E}.  Within one month of the introduction of the WWW interface, the usage of the Abstract Service tripled, and it has continued to rise ever since (\\cite{1997Ap&SS.247..189E}).  With the increased usage of the system due to the easy access through the WWW, severe limitations of the underlying commercial database system very quickly became apparent.  We soon moved to an implementation of the search engine that was custom-built and tailored to the specific requirements of the data that we used.  In January 1995 we started to provide access to scanned journal articles (\\cite{1996adass...5..558A}).  The user interface to these scans provided the user with the capability to access the scans in various formats, both for viewing and for printing.  With time, other interfaces to the abstracts and scanned articles were developed to provide other data systems the means to integrate ADS data into their system (\\cite{1996adass...5..569E}).  With the adoption of the WWW user interface and the development of the custom-built search engine, the current version of the ADS Abstract Service was basically in place.  The following sections describe the current status of the different access capabilities (sections~\\ref{dataretrieval} and \\ref{cookies}), the search engine (sections~\\ref{searchengine} and \\ref{optimization}), access statistics for the ADS system (section~\\ref{accessstats}), and future plans for the ADS interface and search engine (section~\\ref{future}). ",
        "conclusions": "The ADS Abstract Service has been instrumental in changing the way astronomers search and access the literature.  One of the reasons for the immediate acceptance of the ADS by the astronomical community when the WWW based version became available was the ease of use of the interface.  It provided access to many advanced features, while still making it extremely easy to execute simple queries.  Most of the use is for simple queries, but a significant number of queries utilize the more sophisticated capabilities.  The search engine that provides this access was crucial to the success of the ADS as well.  The most important aspects of a search engine are its speed and its flexibility to accommodate special features.  The custom-designed software of the ADS search engine proved that at least in some instances a custom design has considerable advantage over a general purpose system.  The search speed that we were able to achieve and the flexibility of the custom design that allows us to quickly adapt to our users' needs have justified the efforts of developing a system from scratch that is tailored to the specific data set.  The ADS had 31,533 users issue 780,711 queries and retrieve 15,238,646 references, 615,181 abstracts, and 1,119,122 scanned pages in November, 1999.  Since the start of the ADS we have served 311,594.261 references and 17,146,370 scanned article pages.  These numbers speak for them self about the success of the ADS."
    },
    "0002/astro-ph0002428_arXiv.txt": {
        "abstract": "The evolution of rotating, isolated  clusters of stars up to core-collapse is investigated with  n-body numerical codes. The simulations  start off from axisymmetric generalisations of King profiles, with added global angular momentum. In this contribution we report on results obtained  for two sets of single-mass cluster simulations. These confirm the more rapid evolution of even mildly-rotating clusters.  A model is presented with rotational energy comparable to $\\omega$Centauri's; it reaches core-collapse in less than half the time required for  non-rotating clusters. ",
        "introduction": "Star clusters are  % self-gravitating Newtonian systems of choice where to  brew complex gravitational dynamics (Meylan \\& Heggie 1997 for a review). Observations of old, globular, stellar clusters have led to the formulation of spherically symmetric dynamical models of equilibria. The most successful and universally studied one-integral  spherical models are the King (1966) profiles. However, surveys of up to 100  Milky Way clusters have found small but significant departures from spherical symmetry (White \\& Shawl 1987): fits to their projected isophotes yield ellipticities $<\\!\\! \\epsilon\\!\\!> \\equiv 1 - <\\!\\! a/b\\!\\!> \\approx 0.07 \\pm 0.01$.    A study of 173 clusters in M31 found  $<\\!\\!\\epsilon\\!\\!> = 0.09 \\pm 0.04$ (Staneva, Spassova \\& Golev 1996). Observations  of young clusters in the Large Magellanic Cloud revealed  isophotal contours  with ellipticities as large as $ \\epsilon = 0.3$ (Elson, Fall \\& Freeman 1987; Kontizas et al. 1990). This raises  the possibility  that   clusters    are   formed     as  strongly     flattened structures which then  evolve towards rounder configurations (cf. Frenk \\& Fall 1982; Boily, Clarke \\& Murray 1999; Theis \\& Spurzem 1999), and brings up important theoretical issues concerning  processes which may  drive this evolution.  Rotation stretches any stellar association  along a preferred  axis: observations of  the clusters $\\omega$Centauri and M13 have  shown that they are flattened by rotation (Meylan \\& Mayor 1986; Merritt,  Meylan \\& Mayor 1997; Lupton, Gunn \\& Griffin 1987).  Thus angular momentum, measured or possibly % lost during evolution, offers  a way  to account for the morphology of  clusters. Yet to date there are few evolutionary models of clusters with initial angular momentum. We have started on  a project to develop three-dimensional dynamical models of rotating star clusters. In this articles results for two n-body models of isolated clusters are presented.  Previous modelling of rotating clusters is reviewed  first. ",
        "conclusions": "The faster evolution of clusters with rotation has been illustrated with two sample runs. The time to core-collapse we found from three-dimensional n-body simulations are in  agreement with two-dimensional Fokker-Planck calculations (Einsel \\& Spurzem 1999): the collapse time of 5.4 $\\tcol$ obtained for the $\\omega_o = 0.8$ agrees with  the Fokker-Planck solution of 5.6 $\\tcol$ for these parameters. This increases confidence in the results up to core collapse, obtained with two different algorithms."
    },
    "0002/astro-ph0002452_arXiv.txt": {
        "abstract": "Previously we have shown (Maller et al, 1998) that the kinematics of Damped Lyman Alpha Systems (DLAS) as measured by Prochaska and Wolfe (1998) can be reproduced in a multiple disk model (MDM) if the gaseous disks are of sufficient radial extent.  Here we discuss this model's predictions for the relationship between DLAS and Lyman break galaxies (LBGs), which we here take to be objects at $z \\sim 3$ brighter than ${\\mathcal{R}}=25.5$. We expect that future observations of the correlations between DLAS and LBGs will provide a new data set able to discriminate between different theoretical models of the DLAS. Djorgovski (1997) has already detected a few optical counterparts and more studies are underway. ",
        "introduction": " ",
        "conclusions": "We have used the Semi Analytic Models (SAMs) of Somerville \\& Primack (1999) to determine the distribution of galaxies in a dark matter halo, and the amount of cold gas in each galaxy.  The SAMs also contain the star formation history of each galaxy so we can explore the optical properties of the galaxies in the halos that give rise to DLAS.  Here we present results on the optical counterparts from the two models discussed in Maller (1999), in which we matched the kinematic data with thicker and less radially extended, or thinner, more radially extended gas disks.  We show both models here only to demonstrate that the optical properties are not highly sensitive to the details of the gas modeling.  We only refer to optical counterparts that reside in the same virialized halo that produces the DLAS. The contribution of LBGs in neighboring dark matter halos will be explored in future work, but is expected to be relatively unimportant.  The properties of the optical counterparts will place strong constraints on DLAS models.  One constraint is the number of DLAS with optical counterparts.  Figure~\\ref{figopt}a shows the distribution that we see in our models.  Eighty percent of DLAS do not have an optical counterpart with ${\\mathcal{R}} < 25.5$, while a rare five percent contain two or more such galaxies in the same halo.  Lastly we show the distribution of optical impact parameter (Figure~\\ref{figopt}b) in our models.  The optical impact parameter is the physical distance between the line of sight to the quasar and the centroid of the light distribution of the LBG.  We obtain a broad distribution of optical impact parameter values in our model.  Because of the large radial extent of our gas disks, the DLAS are often many stellar disk scale lengths from the center of the light distribution. Also in the MDM scenario, with many galaxies in a single halo, sometimes the galaxy bright enough to be identified as an optical counterpart is not one of the galaxies giving rise to the DLAS: in this case very large separations are possible.  Thus we expect the predictions about the optical impact parameter to be unique to the multiple disk model, and a useful way of distinguishing it from other models.  \\begin{figure} \\plottwo{optdis.ps}{opt_b.ps} \\caption{ a) The distribution of the number of optical counterpart of DLAS. b) The distribution of optical impact parameters.  In both panels the solid line is the thicker gas disk model, while the dashed line is for the thinner gas disks. }\\label{figopt} \\end{figure}"
    },
    "0002/astro-ph0002208_arXiv.txt": {
        "abstract": "We show that the available constraints relating to $^{6}$LiBeB Galactic evolution can be accounted for by the so-called superbubble model, according to which particles are efficiently accelerated inside superbubbles out of a mixture of supernova ejecta and ambient interstellar medium.  The corresponding energy spectrum is required to be flat at low energy (in $E^{-1}$ below 500~MeV/n, say), as expected from Bykov's acceleration mechanism.  The only free parameter is also found to have the value expected from standard SB dynamical evolution models.  Our model predicts a slope 1 (primary) and a slope 2 (secondary) behaviour at respectively low and high metallicity, with all intermediate slopes achieved in the transition region, between $10^{-2}$ and $10^{-1}Z_{\\odot}$. ",
        "introduction": "Galactic nucleosynthesis and chemical evolution are about how a given element is produced in the universe and how its abundance evolved from the primordial universe on.  In the case of the light elements, it is widely agreed that the nucleosynthesis occurs through spallation reactions induced by energetic particles (EPs) interacting with the interstellar medium (ISM).  In these reactions, an heavier nucleus (most significantly C, N or O) is `broken into pieces' and transmuted into one of the lighter $^{6}$Li, $^{7}$Li, $^{9}$Be, $^{10}$B or $^{11}$B nuclei.  Except for $^{7}$Li, this spallative nucleosynthesis is thought to be the main (if not the only) light element production mechanism.  The case of $^{11}$B is slightly more complicated, as neutrino-induced spallation in supernovae (the so-called $\\nu$-process) is sometimes invoked to increase the B/Be and $^{11}\\mathrm{B}/^{10}\\mathrm{B}$ ratios which one would expect should the light elements be produced by nucleo-spallation alone.  Concerning the Galactic evolution of light element abundances, Fields et al.  (2000) have recently re-analyzed the available data as a function of O/H, discussing the uncertainties associated with the methods used to derive the O abundance, the stellar parameters and the incompleteness of the samples.  According to their results, Be and B evolution can be described in terms of two distinct production processes: i) a primary process dominating at low metallicity and leading to a linear increase of the Be and B abundances with respect to O -- `slope 1' -- followed by ii) a secondary process compatible with the standard expectations of the Galactic cosmic ray nucleosynthesis scenario (GCRN) -- `slope 2'.  This behaviour is characterized by a transition metallicity, $Z_{\\mathrm{t}}\\equiv(\\mathrm{O/H})_{\\mathrm{t}}$, below which the Be/O and B/O ratios are constant and above which they are proportional to O/H. Although the value of $Z_{\\mathrm{t}}$ is rather uncertain because very few data points have yet been reported at $Z < Z_{\\mathrm{t}}$, energetics arguments show that a primary process \\emph{is} indeed required below, say, $10^{-2}Z_{\\odot}$ (Parizot \\& Drury 1999a,b,2000b, Ramaty et al.  2000), and therefore the very existence of a transition metallicity separating a primary from a secondary evolution scheme seems reasonably well established.  In spite of the current large uncertainties on the exact value of $Z_{\\mathrm{t}}$, Fields et al.  find a range of possible values between $10^{-1.9}$ and $10^{-1.4}(\\mathrm{O/H})_{\\odot}$ (see also Olive, this conference).  This two-slope picture seems to reconcile the two competing theories for light element nucleosynthesis, namely GCRN which predicts a secondary behaviour for Be and B evolution (e.g. Vangioni-Flam et al. 1990, Fields \\& Olive 1999), and the superbubble model which predicts a primary behaviour at low metallicity (Parizot \\& Drury 1999b,2000b, Ramaty et al.  2000).  However, we show here that when applied to the whole lifetime of the Galaxy (not only to the early Galaxy), the superbubble model \\emph{alone} actually predicts the entire two-slope behaviour inferred from observations, and accounts for all the qualitative and quantitative constraints currently available. Implications for particle acceleration inside a superbubble (SB) are also analyzed. ",
        "conclusions": "The SB model described above has been shown to be fully consistent with the qualitative and quantitative constraints of LiBeB Galactic evolution: 1) it explains the inferred two-slope behaviour in the framework of one sole model; 2) it provides the correct value of Be/O at low metallicity; 3) it predicts the correct value of the transition metallicity; 4) it does not break the Spite plateau; 5) it is consistent with the $^{6}$Li/$^{9}$Be ratio at any metallicity.  Most importantly, these successes rely on the value of only one free parameter, namely the proportion of SN ejecta inside a SB. The value which we find is of the order of a few percent, i.e. exactly in the range derived from standard SB dynamical evolution.  Likewise, the SB model is found to be successful only if the SBEPs have the SB spectrum, i.e a flattened shape at low energy (in $E^{-1}$).  But this is exactly what is predicted by the SB acceleration model of Bykov et al.  As a conclusion, the SB model appears to account for all the available constraints about LiBeB evolution by making only the most standard assumptions about the involved models relating to other fields of astrophysics.  This may be considered as lending support to these models as well."
    },
    "0002/astro-ph0002514_arXiv.txt": {
        "abstract": "Radio astronomy provides a unique window on the universe, allowing us to study: non-thermal processes (galactic nuclei, quasars, pulsars) at the highest angular resolution using VLBI, with low opacity. It is the most interesting wave band for SETI searches. To date it has yielded 3 Nobel prizes (microwave background, pulsars, gravitational radiation). There are both exciting possibilities and substantial challenges for radio astronomy to remain at the cutting edge over the next 3 decades. New instruments like ALMA and the SKA will open up new science if the challenge of dealing human generated interference can be met. We summarise some of the issues and technological developments that will be essential to the future success of radio astronomy. ",
        "introduction": "Moore's law for the growth of computing power with time (ie a doubling every 18 months) is often quoted as being vitally important for the success of the next generation of radio telescopes (working at cm wavelengths) such as the Square Kilometre Array (SKA). It is worth noting that radio astronomy has enjoyed a Moore's law of it own, having a exponentially improvement in sensitivity with time as shown in Figure \\ref{fig-1}. In fact the doubling time is approximately 3 years, and has been in progress since 1940, giving an overall improvement in sensitivity of $10^5$. However, we are approaching the fundamental limits of large mechanical dishes and noise limits of broad band receivers systems. We need to look to other means of extending this growth into the future.  \\begin{figure}[htbp] \\plotone{sensitivity_curve3.ps} \\caption{Exponential growth of Radio Telescope sensitivity. Boxes indicate the sensitivity attained when the systems were first commissioned. Since then substantial improvements have been made to the system temperature and bandwidth of many telescopes. This diagram does not convey anything about their present capabilities. Projected capabilities are shown for ALMA, LOFAR, SKA (with 1,10,100 beams) and Arecibo (with 100 beams). The curent 13 beams system on Parkes is also shown.} \\label{fig-1} \\end{figure}  Why do we want to be on this exponential growth curve ? Fields of research continue to produce scientific advances while they maintain and exponential growth in some fundamentally limiting parameter. For radio astronomy sensitivity is definitely fundamentally limiting. An interesting question is whether there are other parameters for which as exponential growth could be maintained for a period of time. The first and most obvious point about exponential growth is that it cannot be sustained indefinitely.  Can we maintain it for sometime into the future ? There are 2 basic ways to stay on the exponential curve: 1) Spend more money, 2) Take advantage of technological advances in other areas.  \\subsection{International Mega Science Projects}  International cooperation is now needed for dramatic improvements in sensitivity, because no one country can afford to do it alone. ALMA, the Atacama Large MM Array being developed by the USA, Europe and Japan is an example. It will cost around \\$US700M spread over 1999-2007 and will provide an unprecedented opportunity to study redshifted molecular lines. The Square Kilometre Array (SKA) which will work at centimetre wavelengths is likely to be a collaboration of 10 or more nations, spending \\$US500M over 2008 to 2015.  There are 2 basic approaches to funding such large projects: 1) User pays, where member countries pay for slice of the time, or 2) The member countries build the facility and make it openly accessible to all. Optical astronomy has moved very much down path 1) as the Keck, VLT, Gemini type facilities demonstrate. Radio astronomy has traditionally followed path 2) as have projects like CERN. For future facilities there may be pressure to move more into the user pays regime, possibly bringing about substantial change in the dynamics of the radio astronomy community.  In the past considerable extensibility was attained by each country learning from the last one to build a telescope. That evolution is very clear from Figure \\ref{fig-1}. Since we are likely to be moving to internationally funded projects, that path to extensibility is much more restricted and designers must think very carefully about designing extensibility into the next generation telescopes.  \\subsection{Extensibility Through Improved Technologies}  For a given telescope, past and present extensibility has been achieved by improvements in 3 main areas:  \\begin{description}  \\item{\\bf System Temperature:} Reber started out with a 5000 K system temperature. Modern systems now run at around 20 K, meaning that if everything else was kept constant, Reber's telescope would now be 250 times more sensitive than when first built. There are possibilities of some improvements in future, but nothing like what was possible in the past.  \\item{\\bf Band Width:} Telescopes like the GBT (Green bank Telescope) having bandwidths some 500 times greater than Reber's, will give factors of 20--25 improvement in sensitivity. Some future improvements will be possible, but again they will not be as large as in the past.  \\item{\\bf Multiple Beams:} Whether in the focal or aperture plane, multiple beam systems provide an excellent extensibility path, allowing vastly deeper surveys than were possible in the past. Although multiple beam systems have been used for a number of experiments in the past, the full potential of this approach is yet to be exploited. A notable example that has made a stride forward in this direction is the Parkes L band system (Stavely-Smith et al. 1996, PASA, 13, 243). The fully sampled focal plane phased array system being developed at NRAO by Fisher and Bradley highlights the likely path for the future.  \\end{description}  Using these three methods, a small telescope like the Parkes 64m has remained on the exponential curve and the forefront of scientific discovery for 35 years (shown in Figure \\ref{fig-1} in 1962 and in 1997). Other telescopes have of course undergone similar evolution and we only high light Parkes as an example. Scope for continuing this evolution looks good for the next decade, but beyond that more collecting area will be needed. A 64 beam system installed in 2010, would allow Parkes to stay on the curve for some time. The technology to do this is probably only 3-4 years away from having a realisable system, making it possible to jump well ahead of the curve. Putting a 100 beam system on Arecibo by 2005 is possible and would allow Arecibo to jump way out in front of the curve as it did when first built in 1964.  The relevance of Moore's law in this context, is that if it continues to hold true for the next 1-2 decades, it will provide the necessary back end computational power to realise the gains possible with multiple beams. ",
        "conclusions": "The possibilities for the future of radio astronomy are good, but there are some challenges issues for the community to consider and address:  \\noindent $\\bullet$ Whole of radio spectrum needed for redshifted lines\\\\ $\\bullet$ About 2\\% of spectrum is reserved for passive use by regulation so must develop other approaches\\\\ $\\bullet$ We cannot (and don't want to) impede the telecommunications revolution\\\\ $\\bullet$ Radio astronomy has low credibility until we use advanced techniques\\\\ $\\bullet$ Essential to influence government policy\\\\ $\\bullet$ Astronomers should have a uniform position\\\\ $\\bullet$ Threatening language doesn't help\\\\ $\\bullet$ Is interference harmful?\\\\"
    },
    "0002/astro-ph0002065_arXiv.txt": {
        "abstract": "We have detected an X-ray flare on the very low mass star VB 10 (GL 752 B; M8 V) using the ROSAT High Resolution Imager.  VB 10 is the latest type (lowest mass) main sequence star known to exhibit coronal activity. X-rays were detected from the star during a single 1.1-ksec segment of an observation which lasted 22 ksec in total.  The energy released by this flare is on the order of $10^{27}$ ergs s$^{-1}$.  This is at least two orders of magnitude greater than the quiescent X-ray luminosity of VB 10, which has yet to be measured.  This X-ray flare is very similar in nature to the far ultraviolet flare which was observed by Linsky et al. (1995) using the Goddard High Resolution Spectrograph onboard the Hubble Space Telescope.  We discuss reasons for the extreme difference between the flare and quiescent X-ray luminosities, including the possibility that VB 10 has no quiescent ($10^6$ K) coronal plasma at all. ",
        "introduction": "Stars at the very bottom of the main sequence have been a very popular topic of research during the last few years, including in studies of stellar activity. Fleming, Schmitt, \\& Giampapa (1995) demonstrated that fully-convective M dwarfs, i.e. those with masses less than 0.3 M$_\\odot$, exhibit high levels of coronal activity, although at that time, virtually no X-ray detections had been made of stars less massive than VB 8 and LHS 3003 (both type M7.) The star VB 10 (a.k.a. GL 752 B, LHS 474; type M8) appeared to have been marginally detected once by the Einstein High Resolution Imager and was catalogued by Barbera et al. (1993).  Then Linsky et al. (1995) reported the detection of a large amplitude flare on VB 10 in the far UV using the Goddard High Resolution Spectrograph. More recently, Neuhauser \\& Com\\`oron (1998) have detected X-ray emission from a very young brown dwarf in Chameleon.  This confirms that activity does exist for the very lowest mass stellar (or substellar) configurations.  We have used the ROSAT High Resolution Imager (HRI) to reobserve VB 10, which was not detected by the ROSAT Position Sensitive Proportional Counter (PSPC) during either the All-Sky Survey or a 7.3-ksec pointed observation (Fleming et al. 1993).  By obtaining a deeper exposure, we hoped to discover the nature of the VB 10 corona, detect its quiescent X-ray luminosity, and compare its level of coronal activity to that of the more massive M dwarfs.  We have succeded in detecting X-ray emission from VB 10, but only during a brief flare similar to the one observed by Linsky et al. (1995) at ultraviolet wavelengths.  In this paper, we present the latest ROSAT data on VB 10.  Section 2 contains a description of the observations and data analysis.  In Section 3, we discuss a hypothesis to account for the stark contrast between the flare and non-flare X-ray flux values, including the possibility of a total lack of $10^6$ K coronal plasma about VB 10. ",
        "conclusions": "This X-ray flare on VB 10 is reminiscent of the far-UV flare which was detected on VB 10 by Linsky et al. (1995).  These authors observed C II,IV and Si IV emission lines only during the last five minutes of an hour-long exposure taken with the GHRS onboard HST.  They concluded that the flare which they had observed indicated increased magnetic heating rates for low-mass stars near the hydrogen-burning mass limit.  In both the UV and ROSAT X-ray observations, quiescent (i.e. non-flare) emission was never detected from VB 10.  For the UV flare, the emission line fluxes were an order of magnitude greater than the upper limits placed on the non-flare emission by Linsky et al. (1995).  For this most recent X-ray flare, the contrast is even greater.  The peak flare luminosity is {\\it at least} more than two orders of magnitude greater than the non-flare value.  For ease of comparison to more massive M dwarfs, we will normalize L$_X$ by L$_{bol}$, which for VB 10 is $1.74 \\times 10^{30}$ ergs s$^{-1}$ based on the absolute K magnitude measured by Leggett (1992) and the bolometric correction of Veeder (1974). This gives us log (L$_X$/L$_{bol}) > -2.8$ for the peak of the flare and log (L$_X$/L$_{bol}) < -5.0$ outside of flare.  Using the data of Fleming et al. (1995) for the late (later than M5), presumably fully-convective M dwarfs within 7 pc of the Sun, we find that these star have values of log (L$_X$/L$_{bol}$) which are typically $-3.5$.  But during flares (e.g. AZ Cnc; Fleming et al. 1993), these stars can reach values of (L$_X$/L$_{bol}) = -3.0$ to $-2.5$.  Therefore, the magnitude of the flare on VB 10 is completely consistent with that of flares on more massive M dwarfs.  In Fig. 3 we display a plot of normalized X-ray luminosity versus absolute visual magnitude for all known M dwarfs within 7 pc that are later than M5, based on data from a volume-limited survey by Fleming et al. 1995. Inspection of Fig. 3 reveals the sharp contrast between the upper limit to the non-flare X-ray emission in VB 10 and the X-ray emission levels of other late dMe stars. In particular, the upper limit for VB 10 is 1-2 orders of magnitude less than that detected in the late dMe stars. It is comparable to the upper limit for GJ1002 in Fig. 3, a quiescent dM5.5 star that does not exhibit H$\\alpha$ emission nor any reported flare activity. By contrast, VB 10 is a known flare star with (variable) H$\\alpha$ line emission in its spectrum. While one cannot build a theory based on one observation, we cannot help but speculate that our result for VB 10 does indeed reflect a decline in coronal heating efficency near the H-burning mass limit.  But somehow these stars are still able to flare.  We do not understand in detail the mechanisms that give rise to energetic, transient outbursts identified as ``flares'' in the Sun and late-type stars.  However, flares do appear to be the result of instabilities that return stressed systems toward configurations that are characterized by lower potential energy (Rosner \\& Vaiana 1978). Since we do not have an X-ray spectrum, we cannot verify through modeling that the observed emission in VB 10 is consistent with the presence of loop-like magnetic structures.  However, the energetics of the event suggest the possible occurrence of a large volume of flare plasma, implying the presence of large-scale magnetic structures in the atmosphere.  In particular, we can crudely estimate the plausible range of spatial scales that characterize the flare event in the following manner.  In the absence of an actual energy spectrum (such as the pulse-height spectra that were produced by the now-defunct ROSAT PSPC), we assume some plausible flare plasma parameters.  Based on X-ray observations of other flare events on M dwarfs, we adopt a flaring temperature of $T~\\sim~$ 10$^{7}$ K.  We note that a large flare event recorded by the ROSAT PSPC on a late M dwarf star was well-described by a thermal plasma model fit characterized by a temperature of 2 -- 4 $\\times$ 10$^{7}$ K (Sun et al. 1999).  Utilizing the XSPEC analysis package combined with the observed flux from the HRI and the adopted temperature yields a differential emission measure at the flare maximum of EM $> 1.8 \\times 10^{29}$ cm$^{-5}$.  The corresponding volume emission measure is $$ VEM = 4 \\pi R_{*}^2 EM > 1.1  \\times 10^{50}~~cm^{-3}\\,, \\eqno(1) $$ where R$_{*}$ = 0.102 R$_{\\odot}$ (Linsky et al. 1995).  Given this estimate of volume emission measure along with electron densities in the range of $n_e~\\sim$ 10$^{10}$ cm$^{-3}$ to 10$^{11}$ cm$^{-3}$, we find characteristic linear dimensions for the flare in the range of 0.003 R$_{*}$ to 0.30 R$_{*}$. Thus, the flaring plasma covers a large fraction of the stellar surface if $n_e~\\sim$ 10$^{10}$ cm$^{-3}$, but only a small fraction if $n_e~>$ 10$^{11}$ cm$^{-3}$.  We note that the distance traveled by a sound wave in a 10$^{7}$ K plasma in the 3 minute duration of the flare maximum is about 0.7 R$_{*}$.  These estimates, while not conclusive, are consistent with the likely occurrence of large-scale magnetic structures associated with the flare event on VB 10.   The enormous contrast between the X-ray flare luminosity and the upper limit to quiescent emission invites further consideration, especially in view of the likely existence of significant surface magnetic flux and large magnetic structures in VB 10. We note that if VB10 is characterized by quiescent X-ray emission at a level which is typical for the Sun at the maximum in its activity cycle, i.e., log($L_x/L_{bol})~\\approx$ -6.3 (following Schmitt 1997), then we would still not have had sufficient sensitivity to detect it.  The upper limit for non-flare X-ray emission in VB 10 is more comparable to the levels of emission seen in earlier, more massive M dwarfs (non-dMe flare stars) which are themselves characterized by X-ray emission levels that are in excess of solar, or log($L_x/L_{bol})~\\sim$ -6 to -5 (following Fleming, Schmitt \\& Giampapa 1995).  We thus cannot exclude the possibility that VB 10 does indeed have undetected, quiescent X-ray emission at the level of the Sun or that of earlier, quiescent dwarf M stars.  However, inspection of Fig. 3 suggests that the low level of non-flare X-ray emission in VB 10 is unusual with respect to other late-type, active dMe flare stars. We will therefore briefly consider a hypothesis that may account for the extremely low, or even the possible absence of, steady, quiescent heating in this very cool dwarf.  \\begin{figure} \\plotfiddle{f3.ps}{2.0in}{-90}{70}{60}{-250}{350} \\caption{The normalized X-ray luminosity, log$(L_{x}/L_{bol})$, vs $M_{V}$ for M dwarfs later than M5.  Filled circles denote dMe stars.  Upper limits are indicated by arrows.  Flare and non-flare values for certain objects are shown. Data are taken from Fleming et al. (1993) and Giampapa et al. (1996).  Solar values are from Peres et al. (1999). \\label{fig3} }  \\end{figure}  While there is no comprehensive theory for coronal heating---even in the case of the Sun---a common feature of current theories is that the origin of the nonradiative heating of the corona involves the interaction between turbulent convective motions in the photosphere and the footpoints of magnetic loops (Parker 1972, 1983a,b, 1986; van Ballegooijen 1986 and references therein). In this context, our observations of VB 10 may imply that at the photosphere there is simply not enough in the way of random motions at the loop footpoints to jostle the magnetic fields and, in turn, provide sufficient dynamical stresses to lead to detectable plasma heating to coronal ($T~\\sim$ 10$^{6}$ K) temperatures.  For example,  a high magnetic field strength may suppress convective motions, analogous to the situation in sunspots which are also X-ray quiet.  We further note that the very cool and dense photosphere of VB 10 is dominated in content by molecular hydrogen combined with neutral metals bound in molecules.  Thus, there are few ions to couple the magnetic field with the photospheric gas.  We estimate by extrapolation from M dwarf model photospheres (Mould 1976) to the effective temperature of VB 10 (T$_{eff}$ = 2600 K; Linsky et al. 1995) that the ionization fraction, $\\zeta$, in the dense, upper photosphere of VB 10 is $\\zeta~\\sim~$ 10$^{-7}$.  By contrast, $\\zeta~\\sim~$ 10$^{-4}$ in the upper photosphere of the quiet Sun (Vernazza, Avrett, \\& Loeser 1976).  In early M dwarfs we have that $\\zeta~\\sim~$ 10$^{-5}$ (following Mould 1976). Hence, for stars in the temperature-density regime of VB 10, the upper photospheres are characterized by ionization fractions that are 2 orders of magnitude less than that of earlier M dwarfs and 3 orders of magnitude below that of the Sun. Consequently, the interaction between the field and the ambient gas occurs deeper in the photosphere where the ionization fraction is higher but where sufficient energy to produce significant coronal heating is unable to propagate outward.  The consequences for coronal heating of the occurrence of magnetic structures in very cool dwarfs such as VB 10 is summarized in the following argument due to F. Meyer (1999, private communication; see also Meyer \\& Meyer-Hofmeister 1999).  Given that the gas pressure in the corona is typically negligible compared to the magnetic stresses, we must have that the magnetic field configuration, or ``loop'', is force-free, or $\\bf F = (J \\times B)/{\\rm c}$ = 0.  Thus,  the electric currents must flow along the magnetic field lines. Since for steady magnetic phenomena the current density is divergence-free ($\\bf \\nabla \\cdot J$ = 0), currents must also flow through the magnetic footpoints in the photosphere.  In the cool and dense photosphere of stars such as VB 10, the electrical conductivity is so low that any current system rapidly decays.  From this it follows that the magnetic equilibrium must be current-free everywhere. That is, $\\bf J$ = 0 so that $\\bf \\nabla \\times B$ = 0, and ${\\bf B = -\\nabla}{\\rm \\Phi}$.  Given that the potential field is a minimum energy configuration, any further build-up and storage of magnetic field energy is excluded, implying no (or only a relatively low degree) of magnetic field-related heating of the atmosphere.  At the very least, the decay rate of energy at the footpoints must be faster than any energy input derived from the interaction between motions in the upper photosphere and the magnetic loops, effectively quenching any non-radiative heating that might otherwise have occurred.  Clearly, the above picture as outlined cannot explain flare events. Instead, the transient or flare outbursts observed in these very cool stars must arise from more complex magnetic topologies where the storage of magnetic field energy occurs, but which do not have footpoints in the cool, dense underlying photosphere.  In summary, we have confirmed that the M8 dwarf VB 10 (Gl 752 B) does indeed emit X-rays.  It is the lowest mass star on the main sequence which is known to do so.  This emission, however, appears to be only transient in nature. The contrast between the flare X-ray flux and any possible non-flare (i.e., quiescent) X-ray emission is at least two orders of magnitude. We note that X-ray emission in VB 10 at solar levels or even at the levels seen in earlier dwarf M (non-dMe) stars would not be detected at our sensitivity limits.  In comparison to dMe stars later than M5, the upper limit to the non-flare X-ray emission in VB 10 is unusually low.  We have considered a hypothesis that the low ionization fraction and, hence, low conductivity, in the photosphere of VB 10 inhibits coronal heating to the point where quiescent, $10^6$ K coronal plasma may not even exist.  In this scenario, the transient, presumably $10^7$ K coronal plasma which gives rise to the observed X-ray flare would then be associated with topologically complex magnetic field structures that do not have footpoints in the cool photosphere which is, in turn, dominated by neutral atomic and molecular species.  Should this scenario prove to be correct, then there may indeed be a drop in coronal activity at the bottom of the main sequence:  not at the mass where stars become fully convective, as was once suggested, but at the hydrogen-burning mass limit itself."
    },
    "0002/astro-ph0002253_arXiv.txt": {
        "abstract": "We combine new results from the search for microlensing towards the Large Magellanic Cloud (\\lmc) by \\eros2 (Exp\\'erience de Recherche d'Objets Sombres) with limits previously reported by \\eros1 and \\eros2 towards both Magellanic Clouds.  The derived upper limit on the abundance of stellar mass \\macho s rules out such objects as an important component of the Galactic halo if their mass is smaller than  $1 {\\rm M}_{\\odot}$. \\keywords {Galaxy: halo -- Galaxy: kinematics and dynamics -- Galaxy: stellar content -- Magellanic Clouds -- dark matter -- gravitational lensing } ",
        "introduction": "The search for gravitational microlensing in our Galaxy has been going on for a decade, following the proposal to use this effect as a tool to probe the dark matter content of the Galactic halo (\\cite{pac86}).  The first microlensing candidates were reported in 1993, towards the \\lmc\\ (\\cite{aub93}; Alcock et al. 1993) and the Galactic Centre (\\cite{uda93}) by the \\eros , \\macho\\ and \\ogle\\ collaborations.  Because they observed no microlensing candidate with a duration shorter than 10~days, the \\eros1 and \\macho\\ groups were able to exclude the possibility that a substantial fraction of the Galactic dark matter resides in planet-sized objects (\\cite{aub95}; \\cite{alc96}; Renault et al. 1997; \\cite{ren98}; \\cite{alc98}).  However a few events were detected with longer time\\-scales.  From 6-8 candidate events towards the \\lmc , the \\macho\\ group estimated an optical depth of order half that required to account for the dynamical mass of the standard spherical dark halo; the typical Einstein radius crossing time of the events, $t_E$, implied an average mass of about 0.5~M$_\\odot$ for the lenses (Alcock et al. 1997a). Based on two candidates, \\eros1 set an upper limit on the halo mass fraction in objects of similar masses (Ansari et al. 1996), that is below that required to explain the rotation curve of our Galaxy\\footnote{ Assuming the \\eros1 candidates are microlensing events, they would correspond to an optical depth six times lower than that expected from a halo fully comprised of \\macho s.}.  The second phase of the \\eros\\ programme was started in 1996, with a ten-fold increase in the number of monitored stars in the Magellanic Clouds. The analysis of the first two years of data towards the Small Magellanic Cloud (\\smc) allowed the detection of one microlensing event (Palanque-Delabrouille et al. 1998; see also Alcock et al., 1997b). This single event, out of 5.3 million stars, allowed \\eros2 to further constrain the halo composition, excluding in particular that more than 50~\\% of the standard dark halo is made up of $0.01 - 0.5 \\:{\\rm M}_\\odot$ objects (Afonso et al. 1999). In contrast, an optical detection of a halo white dwarf population was reported (Ibata et al. 1999).  In this letter, we describe the analysis of the two-year light curves from 17.5 million \\lmc\\ stars.  We then combine these results with our previous limits, and derive the strongest limit obtained thus far on the amount of stellar mass objects in the Galactic halo. ",
        "conclusions": "After eight years of monitoring the Magellanic Clouds, \\eros\\ has a meager crop of three microlensing candidates towards the \\lmc\\ and one towards the \\smc , whereas 27 events are expected for a spherical halo fully comprised of $0.5 \\:{\\rm M}_\\odot$ objects. These were obtained from four different data sets analysed by four independent, cross-validated programs. So, the small number of observed events is unlikely to be due to bad detection efficiencies.  This allows us to put strong constraints on the fraction of the halo made of objects in the range [$10^{-7}\\:{\\rm M}_\\odot$, $4\\:{\\rm M}_{\\odot}$], excluding in particular at the 95~\\% C.L. that more than 40~\\% of the standard halo be made of objects with up to $1 \\:{\\rm M}_\\odot$.  The preferred value quoted in Alcock et al. (1997a), $f = 0.5$ and $0.5\\:{\\rm M}_\\odot$, is incompatible with the limits in Fig.~\\ref{excl} at the 99.7\\% C.L. (but see the note added below).  What are possible reasons for such a difference? Apart from a potential bias in the detection efficiencies, several differences should be kept in mind while comparing the two experiments. First, \\eros\\ uses less crowded fields than \\macho\\ with the result that blending is relatively unimportant for \\eros . Second, \\eros\\ covers a larger solid angle (43~deg$^2$ in the \\lmc\\ and 10~deg$^2$ in the \\smc ) than \\macho , which monitors primarily the central 11~deg$^2$ of the \\lmc . The \\eros\\ rate should thus be less contaminated by self-lensing that is more common in the central regions~- the importance of self-lensing was first stressed by Wu (1994) and Sahu (1994). Third, the \\macho\\ data have a more frequent time sampling. Finally, while the \\eros\\ limit uses both Clouds, the \\macho\\ result is based only on the \\lmc . For halo lensing, the timescales towards the two Clouds should be nearly identical and the optical depths comparable. In this regard, we remark that the \\smc\\ event is longer than all \\lmc\\ candidates from \\macho\\ and \\eros .  Finally, given the scarcity of our candidates and the possibility that some observed microlenses actually lie in the Magellanic Clouds, \\eros\\ is not willing to quote at present a non zero {\\it lower} limit on the fraction of the Galactic halo comprised of dark compact objects with masses up to a few solar masses.  \\smallskip {\\bf Note added.}  While the writing of this letter was being finalised, the analysis of 5.7 yrs of \\lmc\\ observations by the \\macho\\ group was made public (\\cite{alc00}).  The new favoured estimate of the halo mass fraction in the form of compact objects, $f = 0.20$, is 2.5 times lower than that of Alcock et al. (1997a) and is compatible with the limit presented here. None of the conclusions in this article have to be reconsidered. A detailed comparison of our results with those of Alcock et al. (2000) will be available in our forthcoming publication (\\cite{las00a})."
    },
    "0002/astro-ph0002123_arXiv.txt": {
        "abstract": "The Galactic Bulge has long been assumed to be a largely old stellar population. However, some recent studies based on observations with the HST WF/PC-1 and WFPC2 of stars in the Galactic Bulge have concluded that the old population may not make up more then 30\\% of the total. Other studies using HST/WFPC2 differential studies of `Bulge' globular clusters and field stars have found the bulge to be comparable in age to the Galactic Halo. A complication in all these studies is the presence of a substantial population of stars which mimic a young bulge population, but which may be, and are often assumed to be, foreground disk stars whose reddening and distance distributions happen to mimic a young bulge turnoff. We show, using number counts in HST/WFPC2 colour-magnitude diagrams of both field stars in the Bulge and of two `bulge' and one `disk' globular cluster (NGC6528, NGC6553, and NGC5927) that the stars interpreted as young in fact are foreground disk stars. Thus, we confirm that the bulk of the bulge field stars in Baade's Window are old. The existence of a young {\\sl metal-rich} population cannot, however, be ruled out from our data.  We also test for age and metallicity gradients in the Galactic Bulge between the two low extinction windows Baade's window ($\\ell$=$1\\fdg1, b$=$-4\\fdg8$) and Sagittarius-I ($\\ell$=$1\\fdg3, b$=$-2\\fdg7$). We use the colour-magnitude diagram of a metal-rich globular cluster as an empirical isochrone to derive a metallicity difference of $\\la 0.2$ dex between Baade's window and SGR-I window. This corresponds to a metallicity gradient of $\\la 1.3$ dex/kpc, in agreement with recent near-IR CMD studies. Such a steep gradient, if detected, would require the existence of a short scale length inner component to the Bulge, most likely that prominent in the near infra red, which perhaps forms a separate entity superimposed on the larger, optical Bulge as observed in Baade's window. ",
        "introduction": "The nature and origin(s) of Galactic Bulges are key aspects of any galaxy formation model, and of the Hubble galaxy classification sequence.  However, the present-day properties of bulges in spiral galaxies are not well known (eg Silk \\& Wyse 1993; Wyse, et al. 1997).  Do bulges form late, early or continuously? Are bulges related to halos? To disks? Are they single stellar populations?  The existence of smooth and/or discontinuous gradients in age and/or metallicity in the stellar population(s) in the Galactic Bulge can help to discriminate between these different scenarios, and is the topic of this paper.  Is there evidence for the widely repeated assumption that the Galactic Bulge is old? Recently published colour magnitude diagrams from HST/WFPC2 and HST/WFC1 of the Galactic Bulge and the bulge globular clusters (Vallenari et al. 1996, Ortolani et al. 1995, Holtzman et al 1993) while dominated by fairly old stars, show a substantial population of stars above the dominant old turnoff. Are these foreground disk stars, or is there a minority very young bulge population? Since even a minority young bulge population is of interest, we examine here the limits on young stars in the bulge windows.   One common approach to determine the properties of the Bulge is to study the so called bulge globular clusters (see e.g. Zinn 1996, Ortolani et al. 1995, Minniti 1996). While primarily motivated by observational convenience, the rationale behind this approach is the assumption that these clusters may be valid tracers of the stellar population(s) of the Bulge (see however Zinn 1996 and Harris 1998) Formation scenarios relevant to this approach include the possibility that the Bulge has been assembled from numerous such clusters and these are the last surviving (Gnedin \\& Ostriker 1997), or that the clusters formed a system associated with the Bulge rather than with the rest of the spheroidal component(s) of the Galaxy. Thus the idea is that we may be able to infer the age and/or metallicity of the Bulge stellar population either directly from studies of these clusters, or through differential studies of the clusters and field stars, under the assumption of similar metallicity. Since there is no ab initio understanding of the formation of either galactic bulges or globular clusters, and the age range of the globular cluster system remains a topic of active debate, such analyses merit close scrutiny. The most recent and extensive such analysis is that of Ortolani et al. (1995, 1996) who observed two such globular clusters, NGC6553 and NGC6528, with HST/WFPC2, and deduced that the Bulge has the same age as the Halo.  Analysis of suitably-chosen globular clusters introduces several possible complications.  The first is the major problem of defining a proper population of `bulge' globular clusters. Some of the clusters used, e.g.  Ter7, have recently been shown to be associated with the satellite dwarf galaxy Sgr dSph, rather than with the Galaxy itself.  They may therefore not be representative of the Galactic Bulge.  The method of comparing ridge-lines of globular clusters to infer relative ages requires the clusters in question to have similar metallicities, and relative chemical abundances of the alpha-elements, to avoid an age-metallicity degeneracy (Stetson et al. 1996, VandenBerg et al. 1990, 1996). New results by Cohen et al. (1999) show that NGC6553 may be as much as $\\sim 0.5$ dex more metal-rich than 47 Tuc, illustrating the potentially large effects of metallicity range. The impressive recent study of Rosenberg et al. (1999), indicating a  dispersion in ages for the intermediate metallicity globular clusters, and  a large systematic age difference between the metal-rich and metal-poor clusters illustrates the complexity.  Another potential uncertainty is that the metallicity distribution function of the stars in the galactic Bulge is more similar to that in the solar-neighbourhood (cf. Wyse and Gilmore 1995) than to that for clusters within 5 degrees of the Galactic centre (Minniti 1996; Barbuy et al. 1998). The bulge cluster distribution is both more narrow and less metal-rich than the bulge field stars, complicating any direct comparison.  Clearly, if one wishes to know the age of the bulge field stars, it is desirable to observe the stars in the Bulge directly.  Direct studies of the Bulge are difficult due to the severe crowding towards the central regions of the Galaxy and the large, patchy, reddening along the line of sight.  Several detailed studies of the outer Galactic Bulge exist, providing kinematics and chemical abundance distribution functions (Ibata and Gilmore 1995a, 1995b; Minniti et al. 1995; see also Wyse et al. 1997). For the inner Bulge several analyses of the low reddening Baade's window are available (e.g. Ortolani et al. 1995, Vallenari et al. 1996, Holtzman et al. 1993, Terndrup 1988, Ng et al. 1996), with direct studies of the inner bulge field stars in the near-IR recently also becoming available (Frogel et al. 1999), and even mid-IR ISO photometry (Omont et al. 1999, Glass et al. 1999). Additionally, many recent studies have emphasized the high continuing rate of star formation in the inner bulge/disk. Do these stars diffuse with time the few hundred parsecs into the Sgr and Baade's windows?  The interpretation of extant data is unclear, with a variety of contradictory results.  Vallenari et al. (1996) use a mixture of WF/PC-1 and NTT data while Holtzman et al. (1993) rely exclusively on (the same) WF/PC-1 data.  Both groups reached the conclusion that the Bulge is dominated by a significant young stellar populations. Ortolani et al. (1995) using similar NTT data for Baade's Window found the Bulge to be as old as the Halo.  The confusion among the results from the space based observations should be contrasted with ground based optical observations which find little evidence for a substantial young stellar population(s) in the Galactic Bulge (eg Terndrup 1988), albeit rather far from the centre. Note however that these observations do not cover the main-sequence turn-off and the results are based on the giant branch. The main-sequence turn-off is more sensitive to detection of a significantly younger stellar populations.  Further complication is provided by studies of OH/IR stars, suspected to be of intermediate age, which are common in the inner bulge, or disk (Sevenster et al. 1997).   The distribution function of chemical abundances is a key parameter defining a stellar population. In the outer Galactic Bulge Ibata \\& Gilmore (1995b) derived the relevant distribution function from spectroscopy of K giants. They found a mean abundance of $\\sim -0.2$ dex, with a very wide dispersion. In Baade's window McWilliam \\& Rich (1994) and Sadler et al. (1996) provided similar results.  Sadler et al. (1996) found for 400 K giants a mean abundance of [Fe/H] $=-0.11 \\pm0.04$ dex, with more than half the sample in the range $-0.4 < {\\rm [Fe/H]} < +0.3 {\\rm dex}$.  This is similar to the results from the detailed spectroscopic analysis of McWilliam \\& Rich (1994).  A metallicity gradient has been suspected for fields outside Baade's Window, Terndrup (1988) and Minniti et al. (1995).  The important conclusion from the spectroscopic analyses is that the stellar populations of the inner Galactic Bulge are complex, and that their analysis requires careful consideration of projected disk and other populations.  In this paper we study colour-magnitude diagrams, derived from archival HST/WFPC2 images, for 4 fields and two clusters towards the Galactic Bulge and one ``disk'' cluster. We perform a purely differential study of the properties of the Galactic Bulge population(s), quantifying any systematic offsets and/or gradients in age and/or metallicity in the field population(s).  The paper is organized as follows; in Sect. 2 we detail the observations used and in Sect. 3 describe how we derive the photometric magnitudes from the images. Sect. 4 discusses reddening and distances for the individual fields and clusters, and the utility of the cluster data. Sect. 5 discusses the age of the Bulge, while Sect. 6 asks the question whether a metallicity gradient might be present in the inner Bulge. Sect. 7 includes a summarizing discussion which puts our results into the context of other studies. A brief summary is found in Sect. 8. ",
        "conclusions": "We demonstrate how deep images with high spatial resolutions, only available since the advent of HST, enable us to determine the properties of the Galactic Bulge stellar population(s). In this we have used three measures;  \\begin{itemize}  \\item number counts - to search for a young stellar population  \\item histograms - to search for a metal-rich population   \\item comparison of ridge lines - to derive metallicities and look for internal differences, i.e. gradients in age and metallicity  \\end{itemize}  \\noindent In particular we  find that  \\begin{itemize}  \\item  our results are consistent with no significant young stellar population in the Bulge, some 500pc, or 2-5 scale lengths, from the centre, contrary to some previous studies;  \\item the Galactic Bulge has a mean metallicity equal to that of the old disk;  \\item there is  marginal evidence for a central metallicity gradient;  \\item the colour-magnitude diagram of NGC6553 is complex, requiring care in photometric analyses of this cluster.   \\end{itemize}"
    },
    "0002/astro-ph0002315_arXiv.txt": {
        "abstract": "Smoothing operation to make continuous density field from observed point-like distribution of galaxies is crucially important for topological or morphological analysis of the  large-scale structure, such as, the genus statistics or the area statistics (equivalently the level crossing statistics). It has been pointed out that the adaptive smoothing filters are more efficient tools  to resolve cosmic structures than the traditional spatially fixed filters. We study weakly nonlinear effects caused by two representative adaptive methods often used in smoothed hydrodynamical particle (SPH) simulations. Using framework of  second-order perturbation theory, we calculate the generalized skewness parameters for the adaptive methods in the case of  initially power-law fluctuations. Then we apply the multidimensional Edgeworth expansion method and investigate weakly nonlinear evolution of the genus statistics and the area statistics. Isodensity contour surfaces are often parameterized by the volume fraction of the regions above a given density threshold. We also discuss this parameterization method in perturbative manner. ",
        "introduction": "The large-scale distribution of galaxies is one of the most important sources to study  formation and evolution of  cosmic structures. Now there are two ongoing large-scale redshift surveys of galaxies, the Sloan Digital Sky Survey (SDSS, Gunn \\& Weinberg 1995) and the Anglo-Australian Telescope 2dF Survey (Colless 1998). These surveys will bring us enormous information of three-dimensional galaxy distribution and are expected to revolutionarily improve  our knowledge of the large-scale structure in the universe.  The two-point correlation function of galaxies, or its Fourier transform, the power spectrum  are  simple as well as  powerful tools,  and have been widely used to quantify clustering of galaxies ({\\it e.g.} Totsuji \\& Kihara 1969, Peebles 1974, Davis \\& Peebles 1983). These two quantities are based on  the second-order moments of matter fluctuations. When the fluctuations are Random Gaussian distributed, the two-point correlation function or the power spectrum contains full statistical information of the fluctuations. Even though the initial seed of structure formation is often assumed to be random Gaussian distributed as predicted by the standard inflation scenarios (Guth \\& Pi 1982, Hawking 1982, Starobinsky 1982), this simple assumption has not been observationally established. Moreover, the cosmic structures observed today are more or less affected by nonlinear gravitational evolution. Therefore,    only with these two quantities, we cannot study the large-scale structure properly.   Other statistical methods have  been proposed and is expected to play complimentary roles to the traditional analyses based on the second-order moment.  The skewness parameter characterizes  the asymmetry of one point probability distribution function  (PDF) of the density  field and has been investigated  in deep (Peebles 1980, Fry 1984, Juszkiewicz, Bouchet \\& Colombi 1993, Bernardeau 1994, Scoccimarro 1998, Seto 1999). Beside higher-order moments (such as skewness), there are other statistical approaches designed to directly measure the   geometrical or morphological aspects of galaxy clustering.  Connectivity of the isodensity contour is an interesting target for these approaches.  For example, the  genus statistics were proposed by Gott, Melott \\& Dickinson (1986), the area statistics (equivalently the  level crossing statistics) by Ryden (1988), and percolation analysis by Klypin (1988).   In addition, the Minkowski functionals recently attain much attention ({\\it e.g.} Minkowski 1903, Mecke, Buchert \\& Wagner  1994,  Schmalzing \\& Buchert 1997, Kerscher et al. 1997).  To  analyze observed cosmic structures, smoothing operation becomes crucially important in some cases. The observed galaxies are distributed in point-like manner, but geometrical or morphological analyses, such as, the genus statistics,  are usually  based on continuous (smoothed) density field (see also Babul \\& Starkman 1992, Luo \\& Vishniac 1995). We have traditionally used filters with spatially fixed smoothing radius for analyzing  the large-scale structure. Even though this method is  the simplest  from theoretical point of views, other possibilities  are worth investigated. The local statistical fluctuations due to the discreteness of particles are determined by the number of particles contained in the smoothing kernel. If we use a spatially fixed filter, we can measure smoothed quantities at overdensity regions relatively more accurately than  at underdense regions.  As a result, quality of information becomes inhomogeneous. This inhomogeneity is caused by the simple choice to spatially fix the smoothing radius. There must be more efficient methods to resolve cosmic structure from particle distribution (Hernquist \\& Katz 1989). Actually, Springel et al. (1998) have pointed out that the signal to noise ratio of the genus statistics is considerably improved by using adaptive smoothing methods. Adaptive methods are based on   Lagrangian description,  use nearly same number of ``particles\" (mass elements) to construct smoothed density field (Hernquist \\& Katz 1989) and are expected to be less affected by discreteness of mass elements. Therefore,  it seems reasonable that we can resolve cosmic structures more efficiently, using these methods.  In this article, we perturbatively analyzed   quantitative effects caused by  adaptive smoothing methods. We pay special attention to three representative examples, the skewness parameter, the genus statistics and the area statistics. As the skewness is basically defined by the  one point PDF, we can, in principle, discuss it without making continuous density field. But its analysis is very  instructive to see nonlinear effects accompanied with  adaptive smoothing methods. As a first step, we mainly study  the density field in  real space and do not discuss the  effects of biasing (Kaiser 1984, Bardeen, Bond, Kaiser \\& Szalay 1986, Dekel \\& Lahav 1999).  This article is organized as follows. In \\S 2 we describe basic properties of the adaptive smoothing and introduce its two main approaches. Then perturbative formulas are derived for each of them.  In \\S 3 we discuss the skewness parameters of the density field smoothed by these two adaptive methods. We evaluate them using second-order perturbation theory. Some of  results in this section can be straightforwardly applied to the density field in the  redshift space. Then we discuss weakly nonlinear effects of the  genus and  the area statistics using the multidimensional Edgeworth expansion method explored  by Matsubara (1994). To characterize the isodensity contour, parameterization based on the volume fraction above a given density threshold is often adopted. In \\S 4.1 we discuss this parameterization in perturbative manner. In \\S 4.2 we explicitly evaluate the generalized skewness which is closely related to the genus and the  area statistics. In \\S 4.3 and \\S 4.4, we show the weakly nonlinear effects on these two statistics with various smoothing methods. We make a brief summary in \\S 5. ",
        "conclusions": "Observational analysis of galaxy clustering is one of the central issues in modern cosmology. Various methods have been proposed to quantify the clustering, and many of them ({\\it e.g.} topological analyses of isodensity contour) are based on continuous smoothed density field. However  what we can  observe directly  is distribution of point-like galaxies.  Thus smoothing operation is crucially  important in the  analyses of  the large-scale structure.  From theoretical point of views, filters with spatially constant smoothing radii are  natural choice and have been widely   adopted so far. But we should notice that there are no strong convincing reasons to stick to this traditional  method.  There are  few galaxies in void regions even at  semi-nonlinear scales. In these regions  density field obtained   with fixed smoothing radius might be considerably affected by the discreteness of (point-like)  mass elements,  and might hamper our analyses  of the cosmic structures. Adaptive smoothing is basically Lagrangian description, and we use nearly same number of particles to construct  smoothed density field at each point.   Thus it is quite possible that the  adaptive methods  are more efficient than the  fixed methods  to  analyze the large-scale structure. Actually, Springel et al. (1998) have recently pointed out that using adaptive filters, signal to noise ratio of the genus statistics is improved even at weakly nonlinear scales $\\gsim 10h^{-1}{\\rm Mpc}$.  In this article, we have developed a perturbative analysis of  adaptive smoothing methods  that are applied to quantify the   large-scale structure. Even though adaptive methods might be  promising approaches  in observational cosmology,  this kind of analytic investigation has not been done so far. Our targets are weakly nonlinear effects induced by two typical adaptive approaches, the gather approach and the scatter approach (Hernquist \\& Katz 1989). The concept of these methods is easily understood with equations (8) and (9). The gather approach is  easier to handle analytically.  Numerical costs dealing with discrete particles' systems are also  lower in this approach (Springel et al. 1998). Effects caused by these two  adaptive methods start from second-order of  $\\delta$ in perturbative sense. They   modify quantities which  characterize the nonlinear mode couplings induced by gravity ({\\it e.g.} Peebles 1980).  In \\S 3 we have  investigated  the skewness parameter $S$ which is a fundamental measure to quantify asymmetry  of one point PDF. We have shown that the  skewness for  a gather  top-hat filter  does not depend on the spectral index $n$ in real space, and very weakly depends on it ($S\\simeq 35.2-0.15n$: Einstein de-Sitter background) in redshift space. In the case of Gaussian filter, the skewness parameters show similar behaviors  both in the  scatter and gather approaches. They are increasing functions  of $n$,  in contrast  to the fixed smoothing method.  Next in \\S 4, the genus and area statistics have been studied with Gaussian adaptive filters. Our analysis is based on the multidimensional Edgeworth expansion  explored by Matsubara (1994). We use two quantities $\\nu(\\equiv\\delta/\\sigma)$ and $\\nu_r$ to parameterize isodensity contours. The latter $\\nu_r$ is defined by the volume fraction above  a given density threshold (Gott, Melott, \\& Dickinson 1986). It is explicitly shown that using  this  parameterization, two statistics with the fixed smoothing method are very weakly affected  by  semi-nonlinear gravitational dynamics, as experientally confirmed by N-body simulations. For the  gather smoothing, we found that the $\\nu_r$- parameterization is more effective to keep original linear shape of the area statistics than of the genus statistics.   The parameters $S, T$ and $U$ which characterize the nonlinear corrections of isodensity contour depend largely on the filtering methods. We can characterize nonlinear effects of these methods in somewhat intuitive manner, using results for the genus statistics. The scatter approach makes  low density tails  more meatball-like, but high density tails are transformed in the direction of sponge-like (connected) topology.  This is a remarkable difference from fixed or gather smoothing methods.   Our investigation in this article has been fully analytical one, using perturbative technique of cosmological density field. Numerical analyses based on N-body simulations would play complementary roles to results obtained here,  and thus are  very important. Apart from numerical investigations, perturbative treatment  given in this article would be also developed  further in several ways. The smoothed velocity field is crucially important material in observational cosmology as it is supposed to be less contaminated by effects of biasing (Dekel 1994, Strauss \\& Willick 1995). But our observational information is limited to the line of sight peculiar velocities only at points where astrophysical objects exist. Thus  construction of smoothed velocity field contains similar characters as discussed in this article ({\\it e.g.} Bernardeau \\& van de Weygaert 1996). There is another (more technical) problem that has not mentioned so far. In this article we have only studied spherically symmetric  filter functions. Springel et al. (1998) have shown that signal to noise ratio of the genus curves is further improved by using a triaxial kernel, taking   account of  tensor information of local density field. This point must be  also worth studying."
    },
    "0002/astro-ph0002473_arXiv.txt": {
        "abstract": "We present high--resolution $J$--, $H$--, and $K$--band observations of the carbon star IRC\\,+10\\,216.  The images were reconstructed from 6 m telescope speckle interferograms using the bispectrum speckle interferometry method.  The $H$ and $K$ images consist of several compact components within a 0.2\" radius and a fainter asymmetric nebula. The brightest four components are denoted with A to D in the order of decreasing brightness. A comparison of our images % gives --- almost like a movie of five frames --- insight to the dynamical evolution of the inner nebula. For instance, the separation of the two brightest components A and B increased by almost 40$\\%$ from 191 mas in 1995 to 265 mas in 1998.  At the same time, component B is fading and the components C and D become brighter.  The {\\sf X}--shaped bipolar structure of the nebula implies an asymmetric mass--loss suggesting that IRC\\,+10\\,216 is very advanced in its AGB evolution, shortly before turning into a protoplanetary nebula.  The cometary shape of component A suggests that the core of A is not the central star, but the southern lobe of a bipolar structure.  The position of the central star is probably at or near the position of component B. ",
        "introduction": "IRC\\,+10\\,216 (CW Leo) is the nearest and best--studied carbon star and one of the brightest infrared sources. % It experiences strong mass-loss rates of $\\dot{M} \\approx 2-5\\times10^{-5}$M$_{\\odot}\\,$yr$^{-1}$ (Loup et~al.\\ 1993\\nocite{LoupForveilleEtAl93}).  The central star of IRC\\,+10\\,216 is a long--period variable with a period of $\\sim 649$\\,days (Le~Bertre 1992\\nocite{LeBertre92}). Recent distance estimates range from 110~pc to 150~pc (Groenewegen 1997\\nocite{Groenewegen97}, Crosas \\& Menten 1997\\nocite{CrosasMenten97}). IRC\\,+10\\,216's initial mass can be expected to be close to 4\\,M$_\\odot$ (Guelin et~al.\\ 1995\\nocite{GuelinForestiniEtAl95}, Weigelt et~al.\\ 1998\\nocite{WeigeltBalegaEtAl98}).  The bipolar appearance of the nebula around this object was already reported, e.g., by Kastner \\& Weintraub \\cite{KastnerWeintraub94}.  The non-spherical structure is consistent with the conjecture that IRC\\,+10\\,216 is in a phase immediately before entering the protoplanetary nebula stage.  The most recent high--resolution observations of this object and its circumstellar dust shell were reported by Weigelt et~al. \\cite{WeigeltBalegaEtAl98}, Haniff \\& Buscher \\cite{HaniffBuscher98}, and Tuthill et~al.\\ \\cite{TuthillMonnierEtAl99}.  The results of Dyck et~al.\\ \\cite{DyckBensonEtAl91} and Haniff \\& Buscher \\cite{HaniffBuscher98} showed that the dust-shell structure of IRC\\,+10\\,216 is changing within some years. ",
        "conclusions": "We have presented high-resolution $J$--, $H$--, and $K$--band observations of IRC\\,+10\\,216 with the highest resolution so far at $H$ of 70~mas.  A series of $K$--band images from five epochs between October 1995 and November 1998 shows that the inner nebula is non-stationary. The separations of the four dominant resolved components increased within the 3 years by almost $\\sim40\\%$.  For the two brightest components a relative velocity within the plane of the sky of about 23~mas/yr or 14~km/s was found.  Within these 3 years the rather faint components C and D become brighter whereas component B is fading.  The general geometry of the nebula seems to be bipolar.  We find that the most promising model to explain the structures and changes in the inner nebula is to assume that the star is at the position of component B.  The star then is strongly but not totally obscured at $H$ and $K$.  Consistently component B is very red in the $H-K$ color while A and the northern $J$--band components are relatively blue. Similarly the polarization pattern with strong polarization in the northern arms and still a significant polarization in the peak supports this model. The inner nebula and the apparent motions seem to be rather symmetric around this position and the observed changes are consistent with an enhanced mass loss since 1997.  IRC\\,+10\\,216 is without doubt in a very advanced stage of its AGB evolution. The observed bipolarity of its dust shell even reveals that it has possibly entered the phase of transformation into a protoplanetary nebula."
    },
    "0002/astro-ph0002190_arXiv.txt": {
        "abstract": "The lifetime of the luminous phase of quasars is constrained by current observations to be $10^{6} \\lsim t_Q \\lsim 10^8$ years, but is otherwise unkown.  We model the quasar luminosity function in detail in the optical and X--ray bands using the Press--Schechter formalism, and show that the expected clustering of quasars depends strongly on their assumed lifetime $t_Q$. We quantify this dependence, and find that existing measurements of the correlation length of quasars are consistent with the range $10^{6} \\lsim t_Q \\lsim 10^8$. We then show that future measurements of the power spectrum of quasars out to $z\\sim 3$, from the 2dF or Sloan Digital Sky Survey, can significantly improve this constraint, and in principle allow a precise determination of $t_Q$.  We estimate the systematic errors introduced by uncertainties in the modeling of the quasar-halo relationship, as well as by the possible existence of obscured quasars. ",
        "introduction": "\\label{section:introduction}  A long outstanding problem in cosmology is the synchronized evolution of the quasar population over the redshift range $0\\lsim z\\lsim 5$.  Observations in the optical (Pei 1995) and radio (Shaver et al. 1994) show a pronounced peak in the abundance of bright quasars at $z\\approx 2.5$; recent X--ray observations (Miyaji et al. 2000) confirm the rapid rise from $z=0$ towards $z\\approx 2$, but have not shown evidence for a decline at still higher redshifts. Individual quasars are widely understood to consist of supermassive black holes (BHs) powered by accretion (Lynden-Bell 1967; Rees 1984).  A plausible timescale for quasar activity is then the Eddington time, $4\\times10^7$ ($\\epsilon$/0.1) yr, the e-folding time for the growth of a BH accreting mass at a rate $\\dot M$, while shining at the Eddington luminosity with a radiative efficiency of $L=L_{\\rm Edd}=\\epsilon\\dot M c^2$.  The lifetime $t_Q$ of the luminous phase of quasars can be estimated directly, by considering the space density of quasars and galaxies.  At $z\\sim 2$, the ratio $n_Q/n_G \\sim 3\\times 10^{-3}$ implies the reassuringly close value of $t_Q\\sim t_{\\rm Hub} n_Q/n_G\\sim 10^7$ yr (Blandford 1999 and references therein).  These lifetimes are significantly shorter than the Hubble time, suggesting that the quasar population evolves on cosmic time--scales by some mechanism other than local accretion physics near the BH.  It is tempting to identify quasars with halos condensing in a cold dark matter (CDM) dominated universe, as the halo population naturally evolves on cosmic time--scales (Efstathiou \\& Rees 1988; Haiman \\& Loeb 1998; Kauffmann \\& Haehnelt 2000).  Furthermore, quasars reside in a subset of all galaxies, while the redshift--evolution of the galaxy population as a whole (qualitatively similar to that of bright quasars) has been successfully described by associating galaxies with dark halos (e.g. Lacey \\& Cole 1993; Kauffmann \\& White 1993).  A further link between galaxies and quasars comes from the recent detection, and measurements of the masses of massive BHs at the centers of nearby galaxies (Magorrian et al. 1998; van der Marel 1999).  These arguments suggest that the evolution of the quasar population can indeed be described by ``semi--analytic'' models, associating quasars with dark matter halos. In this type of modeling, the quasar lifetime plays an important role. The quasar phase in a single halo could last longer ($t_{\\rm Q}\\sim10^8$yr), with correspondingly small $M_{\\rm bh}/M_{\\rm halo}$ ratios, or last shorter ($t_{\\rm Q}\\sim10^6$yr), with larger BH formation efficiencies (Haiman \\& Loeb 1998; Haehnelt et al. 1998).  Note that although recent studies have established a correlation between the bulge mass $M_{\\rm bulge}$ and BH mass $M_{\\rm bh}$, this correlation leaves a considerable uncertainty in the relation between $M_{\\rm bh}$ and the mass $M_{\\rm halo}$ of its host halo.  If the initial density fluctuations are Gaussian with a CDM power spectrum, the clustering of collapsed halos is a function of their mass -- rarer, more massive halos cluster more strongly (Kaiser 1984; Mo \\& White 1996).  Hence, measurements of quasar clustering are a potentially useful probe of both BH formation efficiencies and quasar lifetimes (La Franca et al. 1998, Haehnelt et al. 1998).  In this paper, we assess the feasibility of breaking the above degeneracy, and inferring quasar lifetimes, from the statistics of clustering that will be available from the Sloan Digital Sky Survey (SDSS, Gunn \\& Weinberg 1995) and Anglo-Australian Telescope Two-Degree-Field (2dF, Boyle et al. 1999).  Previous works (e.g. Stephens et al.  1997; Sabbey et al. 1999) have yielded estimates suggesting that quasars are clustered more strongly than galaxies. However, the current uncertainties are large, especially at higher redshifts, where clustering has been found to decrease (Iovino \\& Shaver 1988; Iovino et al. 1991), to stay constant (Andreani \\& Cristiani 1992; Croom \\& Shanks 1996), or to increase with redshift (La Franca et al. 1998). As a result, no strong constraints on the life--time can be obtained yet. The key advance of forthcoming surveys over previous efforts is two--fold.  Because of their sheer size, i.e. the large number of quasars covering a large fraction of the sky, both shot--noise and sample variance can be beaten down, significantly reducing the statistical uncertainties.  Furthermore, the large sample-size will eliminate the need to combine data from different surveys with different selection criteria, hence allowing cleaner interpretation.  Recent measurements of the local massive black hole density have stimulated discussions of a radiative efficiency which is much lower than the usual $\\sim 0.1$ (e.g. Haehnelt et al. 1999). A convincing constraint on the lifetime of quasars could be therefore highly interesting, as this might have implications for the local accretion physics near the BH.  This paper is organized as follows. In \\S~\\ref{section:models}, we summarize our models for the quasar luminosity function, and in \\S~\\ref{section:clustering} we compute the quasar correlation function in these models.  In \\S~\\ref{section:present}, we compare the model predictions with presently available data, and in \\S~\\ref{section:future} we assess the ability of future optical redshift surveys to discriminate between the various models. In \\S~\\ref{section:xray}, we repeat our analysis in the soft X--ray band, and examine the contribution of quasars to the X-ray background, and its auto--correlation.  In \\S~\\ref{section:discussion}, we address some of the caveats arising from our assumptions, and in \\S~\\ref{section:conclusions} we summarize our conclusions and the implications of this work. ",
        "conclusions": "\\label{section:conclusions}  In this paper, we have modeled the quasar luminosity function in detail in the optical and X--ray bands using the Press--Schechter formalism.  The lifetime of quasars $t_Q$ enters into our analysis through the duty--cycle of quasars, and we find that matching the observed quasar LF to dark matter halos yields the constraint $10^6 \\lsim t_Q \\lsim 10^8$ yr: smaller lifetimes would imply overly massive BH's, while longer lifetimes would necessitate overly massive halos. This range reassuringly brackets the Eddington timescale of $4\\times10^7$ ($\\epsilon$/0.1) yr.  The main conclusion of this paper is that the lifetime (and hence $\\epsilon$, if the Eddington time is the relevant timescale for quasar activity) can be further constrained within this range using the clustering of quasars: for quasars with a fixed luminosity, longer $t_Q$ implies larger host halo masses, and higher bias.  We find that as a result, the correlation length $r_0$ varies strongly with the assumed lifetime.  Preliminary data from the 2dF survey already sets mild constraints on the lifetime.  Depending on the assumed cosmology, we find $t_Q=10^{7.7\\pm 0.8}$ yr ($\\Lambda$CDM) or $t_Q=10^{8\\pm 0.8}$ yr (OCDM) to within $1\\sigma$ statistical uncertainty. These values are also found to satisfy upper limits on the auto--correlation function of the soft X--ray background.  Forthcoming large quasar samples from SDSS or the complete 2dF survey are ideally suited for a study of quasar clustering, and they can, in principle constrain the quasar lifetime to high accuracy, with small statistical errors. We expect the modeling of the quasar--halo relation, as well as the possible presence of obscured quasars, to be the dominant sources of systematic uncertainty.  Not discussed in depth in this paper is the possibility of using higher moments (such as skewness), which our models also make definite predictions for, and will be considered in a future publication. Remarkably, our best determination of the lifetime of quasars might come from the statistics of high--redshift quasars, rather than the study of individual objects."
    },
    "0002/astro-ph0002159_arXiv.txt": {
        "abstract": "We present VLBA observations of water maser emission associated with the star forming region IRAS 21391+5802, which is embedded in a bright rimmed cometary globule in IC1396. The angular resolution of the maps is $\\sim$ 0.8 mas, corresponding to a spatial resolution of $\\sim$0.6 AU, at an estimated distance of 750 pc. Proper motions are derived for 10 maser features identified consistently over three epochs, which were separated by intervals of about one month. The masers appear in four groups, which are aligned linearly on the sky, roughly along a northeast--southwest direction, with a total separation of $\\sim$520 AU ($\\sim$0$\\rlap{.}''$7). The 3D velocities of the masers have a maximum value of $\\sim$42 km~s$^{-1}$ ($\\sim$9 AU yr$^{-1}$). The average error on the derived proper motions is $\\sim$4 km~s$^{-1}$. The overall pattern of proper motions is indicative of a bipolar outflow. Proper motions of the masers in a central cluster, with a projected extent of $\\sim$ 20 AU, show systematic deviations from a radial outflow. However, we find no evidence of Keplerian rotation, as has been claimed elsewhere. A nearly circular loop of masers lies near the middle of the cluster.  The radius of this loop is 1 AU and the line-of-sight velocities of the masers in the loop are within 2 km~s$^{-1}$ of the systemic velocity of the region. These masers presumably exist at the radial distance where significant dust condensation occurs in the outflow emanating from the star. ",
        "introduction": "An important phase in the evolution of young stellar objects (YSO) is the one during which they exhibit energetic mass outflows, and simultaneous accretion of material onto the star-forming core. Since the discovery of bipolar outflows (Snell et al. 1980), much progress has been made in the study of the large scale and phenomenological characteristics of such outflows (e.g. Bachiller 1996), but relatively fewer studies have  probed the regions within $\\sim$ 100 AU of the YSOs,  where the winds that drive the outflows, originate (e.g., Shu et al. 1994).  A significant fraction of all YSOs exhibit water maser emission, and the maser luminosity appears to be correlated with the mechanical luminosity of the outflows (Felli et al. 1992). The maser emission is thought to arise from dust-laden shocked gas close to the exciting YSOs, where densities are $\\sim$ 10$^{9}$ cm$^{-3}$ and temperatures are $\\sim$400 K (e.g. Elitzur et al. 1989). Since such conditions are likely to be highly localized, and since the velocity coherence required for building up large path lengths for the maser amplification is also likely to occur over relatively smaller regions, the water maser emission provides an excellent tool for probing the kinematics of the gas with high angular resolution ($\\sim$1 mas). Radio interferometric observations of the water maser emission and  measurements of maser proper motions have been carried out mostly towards high-mass star forming regions such as Orion-KL, W49N, W51 Main,  W51 North, Sgr B2 and W3~(OH) (Genzel et al. 1981a;  Gwinn et al. 1992; Genzel et al. 1981b; Schneps et al. 1981; Reid et al. 1988; Alcolea et al. 1993). Similar studies of low and intermediate-mass ($<$10 M$_{\\odot}$) star forming regions with very high spatial resolution are only now becoming available (Claussen et al. 1998, Furuya et al. 1999, Wootten et al. 1999).  IRAS 21391+5802 is an infrared source embedded in a bright-rimmed globule appearing on the northern periphery of the giant H {\\sc ii} region IC1396. Several observational studies of this object at various wavelengths suggest that it is an intermediate-mass star forming region (Sugitani et al. 1989, Wilking et al. 1993, Patel et al. 1995, Cesaroni et al. 1999). Its far-infrared luminosity is 235 L${_\\odot}$ (Saraceno et al. 1996). Water maser emission from this source  was first detected by Felli et al. (1992) and interferometric observations were made by Tofani et al. (1995) using the NRAO\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under a cooperative agreement with Associated Universities, Inc.} Very Large Array (VLA) at a resolution of 0$\\rlap{.}''$1 resolution. These observations revealed a linear distribution of masers, with a systematic velocity gradient along its length. More recently, Slysh et al. (1999) reported single-epoch observations of these masers with mas angular resolution. However, interpretation of the kinematics has remained ambiguous without proper motion data.  We present here the results of a multi-epoch program of interferometric observations of the water maser emission from IRAS21391+5802. In the next section, we summarize the observational details, followed by results and our interpretation using a simple model of a bipolar outflow. ",
        "conclusions": "We have detected proper motion of 10 water maser features associated with the star-forming region IRAS 21391+5802. The water masers trace a bipolar outflow from an intermediate mass  young stellar object. We estimate position and inclination angles of $\\sim 70^{\\circ}$ for the outflow axis, though the opening angle may exceed 40$^{\\circ}$. The maximum observed 3D space velocity of the masers is 42 km~s$^{-1}$. The observed 3D space velocities of the masers suggest the presence of acceleration within the outflow, but we cannot rule out episodic ballistic ejection. Near the center of the flow, there is a roughly circular loop of masers with a radius of $\\sim$1 AU. The masers in this loop are most likely to be tangentially amplified within a shell of dense gas surrounding the YSO. Because there is no evidence for rotation in the source (as in a disk), the shell probably comprises wind material alone. However, dynamics of the IRAS 21391+5802 region are poorly sampled, especially in the loop. Additional observations may still reveal some rotation, or even infall."
    },
    "0002/astro-ph0002229_arXiv.txt": {
        "abstract": "We report the X/$\\gamma$-ray spectrum of GRB971214 and of its afterglow. The afterglow was measured few hours after the main event and for an elapsed time of more than two days. The measure of this GRB and afterglow is relevant due to its extreme, cosmological distance (z=3.42). The prompt event shows a hard photon spectrum, consistent with a broken power law with photon indices $\\Gamma_{\\rm X}\\approx$0.1 below $\\sim$20 keV and $\\Gamma_\\gamma\\approx$1.3 above 60 keV. The afterglow spectrum, measured with the MECS and LECS BeppoSAX telescopes, is consistent with a power law with spectral photon index $\\Gamma$=1.6. Within the statistical accuracy of our measure no spectral evolution is detected during the observation of the afterglow. When integrated during the time span covered by BeppoSAX observations, the power in the afterglow emission, even with very conservative assumptions, is at least comparable with the power in the main event. The IR-to-X rays broad band spectrum is also presented, collecting data from the literature and adding them to the BeppoSAX measure. It shows that the predictions from synchrotron emission models is qualitatively confirmed. The BeppoSAX measurement of the X and $\\gamma$ ray spectrum of this GRB/afterglow is discussed in the framework of current theoretical models ",
        "introduction": "The discovery of X-ray afterglows from Gamma-ray Bursts (GRBs) (Costa et al. \\cite{costa}) is a major step forward to understand this still mysterious phenomenon. The detection of the faint, fading X-ray counterparts of GRBs poses tight constraints to the models for the emission. Multiwavelength studies discovered optical, IR and radio transients associated with the X-ray afterglow, thanks to the unprecedented positioning accuracy obtained with BeppoSAX. The discovery of a substantial redshift (Kulkarni et al. \\cite{kulkarni_c,kulkarni_n}) in the absorption and emission lines in the spectra of the host galaxies associated with the GRBs optical transients puts these catastrophic events at a cosmological distance and results in an extreme energy output from each GRB, if the emission is isotropic.  The recent advances in our knowledge about the cosmic events known as GRBs were mainly due to the accurate positioning allowed by BeppoSAX (Boella et al. \\cite{boella}). This satellite carries on board an optimal set of instruments to detect GRBs (the Gamma Ray Burst Monitor - GRBM Frontera et al. \\cite{frontera}, Feroci et al.  \\cite{ferocigrbm} ), to position them within a few arcminutes (the Wide Field Cameras - WFC; Jager et al. \\cite{jager}) and finally to pinpoint the positions down to tens of arcseconds thanks to rapid (few hours) follow-up observations with the Narrow Field Instruments (Low Energy Concentrator Spectrometer - LECS; Parmar et al. \\cite{parmar}; Medium Energy Concentrator Spectrometer - MECS; Boella et al. \\cite{boella2}; High Pressure Gas Scintillation Proportional Counter - HPGSPC; Manzo et al. \\cite{manzo}; Phoswich Detection System - PDS; Frontera et al. \\cite{frontera} )  The positions given by BeppoSAX (e.g. Piro et al. \\cite{piro} for \\object{GRB960720}) allowed prompt ground based observations with telescopes in optical, radio, IR. Up to now thirteen Optical Transients (OT) were discovered in the error boxes of the X-ray afterglows (van Paradijs et al.  \\cite{ot_1}, Bond \\cite{ot_2}, Halpern et al.  \\cite{halpern}, Groot et al. \\cite{ot_3}, Palazzi et al. \\cite{ot_4}, Galama et al. \\cite{ot_5}, Jaunsen et al. \\cite{ot_6}, Hjorth et al. \\cite{ot_7}, Bloom et al. \\cite{ot_8}, Kulkarni et al. \\cite{ot_9}, Galama et al. \\cite{ot_10}, Palazzi et al. \\cite{ot_11}, Bakos et al. \\cite{ot_12}). After the fading of the OT in most cases a faint galaxy was detected. The detection of the putative host galaxy of GRB971214 is particularly intriguing as the estimate of the redshift is z=3.42, locating this event at an extreme cosmological distance (Kulkarni et al. \\cite{kulkarni_n}).  We observed with BeppoSAX a GRB on December 14.97272 UT 1997 (Heise et al. \\cite{heise_grb}). A follow-up pointing performed 6.67 hours after the main event detected the faint and fading X-ray source \\object{1SAX J1156.4+6513} (Antonelli et al. \\cite{antonelli}). After the fading of the optical transient, spectroscopic observations of the associated host galaxy tentatively put it at cosmological distance, as the estimate of its redshift is z=3.42 (Kulkarni et al. \\cite{kulkarni_n}), that corresponds to a luminosity distance $>$30 Gpc (for H$_0$=65 km s$^{-1}$ Mpc$^{-1}$ and $\\Omega_0$=0.2). The complete evolutionary history of the emitted spectrum from X to $\\gamma$ rays up to 2.5 days after the main event suggests that the afterglow in X-rays begins immediately. The energy output in the afterglow results to be comparable to that in the prompt event. With the new wealth of data on optical counterparts and X--ray afterglows, a major revamping of theoretical models is occurring. With the extragalactic origin firmly established on the basis of observations of host galaxies (Metzger et al. \\cite{metzger}, Kulkarni et al. \\cite{kulkarni_n, ot_9}), the cosmological fireball model (e.g. Cavallo \\& Rees \\cite{cavallo}, M\\'esz\\'aros \\& Rees \\cite{meszrees1}) gives predictions that reasonably fit to observational data. In this paper we discuss the details of the prompt and delayed emission from GRB971214, with emphasis on the X and $\\gamma$ ray spectrum and on its evolution with time. Implications on the models of the prompt event and of the afterglow are discussed. ",
        "conclusions": "Various authors have discussed the emission in the afterglow (e.g. Tavani \\cite{tavani}, Vietri \\cite{vietri}, M\\'esz\\'aros \\& Rees \\cite{meszrees}, Sari et al. \\cite{sari}, Waxman \\cite{waxman_1,waxman_2}), giving a convincing answer to the problem of light curve modeling.\\\\ The observed decay during the afterglow (Heise et al. in preparation) smoothly reconnects with the observed X-ray flux during the burst. In the hypothesis that the afterglow starts a few seconds after the end of the main event, we can estimate the total fluence in the afterglow, to be compared with the X-ray luminosity during the prompt event. This hypothesis is supported by the recent detection of an optical afterglow in GRB990123 a few seconds after the burst trigger (Akerlof et al. \\cite{rotse}) and by the interpretation of the observed complex light curve (Sari \\& Piran \\cite{sapir_grb99}). Also the detection of an early power--law--like tail from GRB920723 (Burenin et al. \\cite{granat}) supports the hypothesis that the afterglow starts immediately after the GRB, and probably without any interruption in the X--ray flux. Of course this assumption adds further uncertainty to the total estimate. As an example, assuming that the afterglow starts $\\sim$1000 s after the trigger of the main event (and not a few seconds after {\\it the end} of the main event) the total integral differs by a factor $\\sim$2.   In doing this estimate we have to assume a spectral shape and slope. Assuming a time decay law $\\propto {\\rm t}^{-\\beta_{\\rm t\\_X}}$ with $\\beta_{\\rm t\\_X}\\sim$1.2 (Heise et al. in preparation), we obtain that the total fluence in the afterglow between 2 and 10 keV uncorrected for redshift is 4.1 $\\times 10^{-7}$ erg \\rcm2. It corresponds to a luminosity of 1.3$\\times 10^{52}$ erg. This is 2 times the total luminosity in the same energy band during the burst. This is a lower limit, as we do not add any time evolution of the spectral shape, but we conservatively adopt the measured spectral index during the afterglow, much steeper than that measured during the prompt event.  If we assume a similar, or even steeper, spectral slope in the 2-700 keV energy interval (uncorrected for redshift), the total energy output from the afterglow is substantial. For the same power law index measured in the 2-10 keV energy band, we obtain a fluence in the afterglow $~4.5\\times 10^{-6} \\flux$. Using a more conservative assumption, adding an exponential cutoff with folding energy $\\sim$50 keV, the total fluence in the afterglow up to 2.5 days after the main event is 1.15$\\times 10^{-6} \\flux$. We conclude that the {\\it radiated} X and $\\gamma$ total luminosity in the afterglow may be estimated to be between 3.6$\\times 10^{52}$ ergs and 1.4$\\times 10^{53}$ ergs. This is comparable to the total radiated power in the main GRB event. This estimate is only speculative, as the spectral data of the afterglow are consistent with any cutoff energy above $\\sim$2 keV (90\\% confidence) but it points out the need for prompt measurements of the afterglow that can give a good estimate of the spectral shape above 10 keV.\\\\ Unfortunately the spectral evolution in hard X-rays (above 10 keV) of most GRBs seems unaccessible to the present generation of telescopes operating in this energy band. This is a very important observational point that cannot be fulfilled up to the next generation of focusing hard X-ray telescopes, maybe at least for a decade.  The time resolved spectral analysis of the afterglow reported in Table 1, even if of low statistical quality, do suggest that the spectral shape remains stable during our observation of the afterglow. While these data cannot be profitably used to constrain the spectral evolution of the afterglow in X--rays, they confirm the stability of the {\\it engine } that is producing the observed X--ray emission in spite of the substantial power--law decline with time in the observed X--ray luminosity.  The spectral evolution from GRB to the X-ray afterglow indicates that the energy distribution shifts towards lower energies, as expected from theoretical models. The model for spectral evolution of Sari et al. (\\cite{sari}) suggests that the high energy tail of the emitted photon spectrum of the afterglow has a power law slope $\\sim-\\frac{\\rm p}{2} - 1$ in the case of fast cooling, where p is the index of a power law distribution of the electrons. Our measurement of a power law slope $\\sim -1.6$ implies p$\\sim 1.2$, a value that would give a non-finite energy in the electrons. In this case the observations can be reconciled with theory assuming that a suitable cutoff, e.g. exponential, in the energy distribution of the parent electron population is present. The observation of this cutoff in the produced photon distribution is however beyond the capability of the present generation of X--ray telescopes and may be accessible to the new missions like {\\it XMM} and {\\it Chandra} if it is substantially below 100 keV. Our data support (as suggested also by Waxman \\cite{waxman_1,waxman_2}) that we observe a regime where the synchrotron cooling time is long compared to the dynamical time. In this regime $\\Gamma = -\\frac{({\\rm p}-1)}{2} - 1$, and for a finite power in the electrons one obtains $\\Gamma < -1.5$, definitely compatible with our measurement. A caveat must be added to this interpretation: as the adiabatic losses become dominant in this regime, the observed {\\it radiated} luminosity is produced via synchrotron losses in much an inefficient way. This brings down the efficiency of the radiative process and in parallel rises accordingly the total energy budget in the afterglow.  The slope of the NIR-to-X-ray afterglow spectrum, corrected for Galactic and local extinction, and its temporal evolution are in fair agreement with models of expanding fireballs.  Our assumption on the intrinsic extinction reasonably conforms with the proposal that GRBs are connected with star formation, and therefore expected to reside in star forming regions of their host galaxies (though not necessarily in extreme starburst galaxies, see Odewahn et al. \\cite{odewahn}). We note that the local extinction correction we adopted has the advantage, with respect to other approaches, of using a specific model curve and of making the spectrum consistent with a single radiation component over more than three decades of frequency.  If the emission is isotropic the total X-$\\gamma$ luminosity in GRB971214, if at redshift z=3.42, is quite higher than 10$^{51}$ ergs, as already pointed out by Kulkarni et al. (\\cite{kulkarni_n}). Our analysis shows that a substantial fraction, more than 60\\%, of the total radiated energy in 2-10 keV is in the afterglow. In addition a reasonable guess of a high energy exponential cutoff, with a folding energy of 50 keV, brings us to conclude that the total power in GRB971214 may be grossly underestimated if based only on the prompt event.  If the efficiency to convert the total energy output from the GRB in the afterglow is low, e.g. $\\sim$10\\%, as usually assumed in most theoretical models of fireball expansion, our measure of a substantial energy output in 2--10 keV during afterglow shows that the total energy balance of a GRB is grossly underestimated. Furthermore, if we consider the probable presence of a high energy tail of the afterglow, a conservative estimate may bring the total energy output from GRB971214 to more than 10$^{54}$ erg for an isotropically emitting source. Alternatively, one may more comfortably stay with a luminosity of 10$^{53}$ erg assuming a more efficient mechanism to effectively extract radiative power from the expanding fireball or assuming an extreme cutoff to the high energy spectrum.  A way out of this deadlock may be beaming (Yi \\cite{yi}, Shaviv \\& Dar \\cite{shdar}). Different authors have discussed the ``beaming solution'' to the GRB/afterglow observational problem (Dar \\cite{dar1,dar2}, Rhoads \\cite{rhoads}, M\\'esz\\'aros \\& Rees \\cite{meszrees}, Drodzova \\& Panchenko \\cite{drodzova}, Panaitescu et al. \\cite{panait}). If the emitted power is strongly beamed, and therefore not isotropically distributed, the total power in the GRB may be reduced by orders of magnitude, depending on the beam open angle. Of course this has some major and obvious impacts. The number of GRBs (not detected at earth) rises by the same orders of magnitude. Limb darkening (due to the random distribution of viewing angles inside the emission cone), time dependent (due to different beaming at different times after the main event) and energy dependent (due to the time--dependent photon energy distribution) effects should be observable once the afterglow sample is large enough. Examples of such effects are discussed in Panaitescu et al. (\\cite{panait}), including a mixed case with a collimated jet and a contribution from isotropic ejecta.  Assuming a jetlike outflow, the models (Panaitescu et al. \\cite{panait}, Rhoads \\cite{rhoads2}, Sari et al. \\cite{sapir_jets}) predict a steeper decay in the light curve, depending on the jet opening angle. This steepening compared to an isotropic fireball expansion occurs for a 10$^{\\rm o}$ opening angle at approximately 6 days after the event, earlier for smaller angles. We do not detect such a steepening in the X--ray light curve up to 2.5 days after the main event. If we assume a beam open angle $\\theta \\geq 10^{\\rm o}$ the total energy in the event is reduced accordingly, compared to the isotropic case.  An argument to assess jetlike or spherical emission is proposed by Sari et al. (\\cite{sapir_jets} ), using the decay slope of the high energy afterglow. As reported by other authors (see above), they suggest that the power law decay for a jetlike emission is appreciably steeper. For an isotropic fireball the expected decay is t$^{-\\beta_{\\rm X}}$ with $\\beta_{\\rm X} \\sim 1.1-1.3$, while for an expanding jet the decay follows a power law with $\\beta_{\\rm X} \\sim$ 2.4 .\\\\ This effect is purely geometrical, as it is geometrical the effect of the jet {\\it ``spillover'' } (Rhoads \\cite{rhoads2}) expanding sideways at the local sound speed. In order to maintain the observed time decay power law $\\propto t^{-1.2}$, the Lorentz factor $\\gamma$ must be $> \\theta^{-1}$ during all the afterglow observation (see e.g. Piran \\cite{piran}). In the case of GRB971214, this must hold up to the last observation of the power law decay, performed approximately 2.5 days (60 hours) after the prompt event. Following Sari et al. (\\cite{sapir_jets}) this translates in a lower limit to the beam opening angle $\\theta_0 > 0.1 \\times (6.2 \\times 60 \\times ({\\rm E}_{52}/ {\\rm n}_1)^\\frac{1}{3} )^\\frac{3}{8} \\approx 0.2$, assuming the total ``isotropic'' energy in the afterglow is $\\sim 10^{53}$ ergs. This lower limit in beaming angle translates into a lower limit in the total radiated power from this GRB (prompt+afterglow) $\\sim 10^{52}$ ergs.  If the effect of beaming is comparable during the prompt event and the afterglow, the conclusions we draw on the relative {\\it observed} energy output apply also directly to the total energy balance in the two cases. If the effect of beaming is evolving from an extremely beamed emission during the event to a relatively hollow beaming during the part of the afterglow we observed, the relative energy balance between the two cases should scale accordingly. As a consequence, the energy budget in the afterglow may increase substantially with respect to that in the prompt event.  In conclusion the measurement of the spectrum during the prompt event and during the afterglow strongly supports the models for synchrotron emission from GRB afterglows, with an agreement both in the X-ray band and in a broader NIR-to-X-rays band (see Fig. 4). The measured spectral slope is in fair agreement with the requirements of the models of expanding fireballs, considering also the observed temporal decay. A simple argument, based on recent models on jetlike emission and on the expected temporal decay in this case, supports the observation of a spherical expanding shell or of a moderate beaming ($\\theta > 0.2$). If this is the case, the observed radiated power in the afterglow is substantial and should be accordingly reproduced in any model for X-ray afterglow from GRBs.  {\\bf Acknowledgements}. This research is supported by the Agenzia Spaziale Italiana (ASI) and the Consiglio Nazionale delle Ricerche (CNR) of Italy. BeppoSAX is a joint program of ASI and of the Netherlands Agency for Aerospace Programs (NIVR). We wish to thank all the people of the BeppoSAX Scientific Operation Centre and the Operation Control Centre for their skillful and enthusiastic contribution to the GRB research program."
    },
    "0002/hep-ph0002236_arXiv.txt": {
        "abstract": "We study a double inflation model (a hybrid inflation $+$ a new inflation) in supergravity and discuss the formation of primordial black holes (PBHs) with mass $\\sim 10^{-20}-10^{5}M_{\\odot}$. We find that in a wide range of parameter space, we obtain PBHs which amount to $\\Omega \\simeq 1$, i.e., PBH dark matter. Also, we find a set of inflation parameters which produces PBHs evaporating now.  Those PBHs may be responsible for antiproton fluxes observed by the BESS experiment. ",
        "introduction": "In the framework of supergravity the reheating temperature of inflation should be low enough to avoid overproduction of gravitinos~\\cite{Weinberg,Ellis}. The new inflation model~\\cite{New-inflation} generally predicts a very low reheating temperature and hence it is the most attractive among many inflation models~\\cite{Linde-text}. However, the new inflation suffers from a fine-tuning problem about the initial condition; i.e., for a successful new inflation, the initial value of the inflaton should be very close to the local maximum of the potential in a large region whose size is much longer than the horizon of the universe.  A framework of a double inflation~\\cite{dynamical-tuning} was proposed to solve the initial value problem of the new inflation model\\footnote{Different models of double inflation were studied by various authors~\\cite{doubleinf}.  }. It was shown that the above serious problem is naturally solved by supergravity effects if there exists a preinflation (e.g., hybrid inflation~\\cite{hybrid-original}) with a sufficiently large Hubble parameter before the new inflation~\\cite{dynamical-tuning}.   In this double inflation model, density fluctuations produced by both inflations are cosmologically relevant if the $e$-fold number of the new inflation is smaller than $\\sim 60$ (the total $e$-fold number $\\sim 60$ is required to solve flatness and horizon problems in the standard big bang cosmology~\\cite{Kolb-Turner}). In this case, the preinflation should account for the density fluctuations on large cosmological scales [including the Cosmic Background Explorer (COBE) scales] while the new inflation produces density fluctuations on smaller scales.  Although the amplitude of the fluctuations on large scales should be normalized to the COBE data~\\cite{COBE}, fluctuations on small scales are free from the COBE normalization and can have arbitrary power matched to observations.  In Ref.\\cite{KKSY-LSS}, a cosmological implication of the double inflation for the large-scale structure formation was discussed. In this paper, we study primordial black hole (PBH) formation in the double inflation model. In Refs.~\\cite{KSY-PBH,KY-PBH}, the production of black hole MACHO was investigated in the double inflation model for a special case\\footnote{Different models for the PBH formation have been studied in Ref.~\\cite{otherPBH}.}. Here, we consider a wide range of parameter space where PBHs are formed\\footnote{In this paper we investigate the PBHs with mass $\\sim 10^{-20} - 10^{5} M_{\\odot}$. The upper bound of $10^5 M_{\\odot}$ is not rigorous but much heavier black holes can be produced in the present model if we take the appropriate model parameters. However, the mass of the black holes should be less than the galactic mass $\\sim 10^{12} M_{\\odot}$, otherwise the power spectrum conflicts the observations ( e.g. distribution of galaxies ).  The lower bound $\\sim 10^{-20} M_{\\odot}$ comes from the requirement that $e$-fold number of the new inflation should be larger than $0$.}. In particular, we show that the double inflation creates small PBHs evaporating now if those PBHs are produced during matter-dominated (MD) era, i.e., before the end of reheating process after the new inflation. We stress that these evaporating PBHs may account for antiproton fluxes observed by the BESS experiment~\\cite{BESS}.  Throughout this paper the gravitational scale ($\\simeq 2.4 \\times 10^{18}$ GeV) is taken to be unity. ",
        "conclusions": "In this paper we have studied the formation of PBHs by taking a double inflation model in supergravity.  We have shown that in a wide range of parameter space PBHs are produced of various masses. These PBHs are interesting since, for example, they may be identified with MACHOs ($M\\sim M_{\\odot}$) in the halo of our galaxy.  Or, they may be PBHs which are just evaporating now ($M\\sim 10^{-19}M_{\\odot}$). Such black holes are one of the interesting candidates for the sources of antiproton fluxes recently observed in the BESS detector~\\cite{BESS}.  The dark matter PBHs play a role of the cold dark matter on the large scale structure formation.  The scales of the fluctuations for PBH formation themselves are much smaller than the galactic scale and thus we cannot see any signals for such fluctuations in $\\delta T/T$ measurements.  However, the PBHs may be a source of gravitational waves.  If the PBHs dominate dark matter of the present universe, some of them likely form binaries.  Such binary black holes coalesce and produce significant gravitational waves~\\cite{Nakamura} which may be observable in future detectors."
    },
    "0002/physics0002036_arXiv.txt": {
        "abstract": "Absolute dielectronic recombination (DR) rates for lithium-like Ni$^{25+}$($1s^2 2s$) ions were measured at high-energy resolution at the Heidelberg heavy-ion storage ring TSR. We studied the center-of-mass energy range 0--130 eV which covers all $\\Delta n$=0 core excitations. The influence of external crossed electric (0--300 V/cm) and magnetic (41.8--80.1 mT) fields was investigated. For the measurement at near-zero electric field resonance energies and strengths are given for Rydberg levels up to n$=$32; also Maxwellian plasma rate coefficients for the $\\Delta n$=0 DR at electron temperatures between 0.5 and 200 eV are provided. For increasing electric field strength we find that for both the $2p_{1/2}$ and the $2p_{3/2}$ series of Ni$^{24+}$($1s^2 2p_j n\\ell$) Ryd\\-berg resonances with $n>30$ the DR rate coefficient increases approximately linearly by up to a factor of 1.5. The relative increase due to the applied electric field for Ni$^{25+}$ is remarkably lower than that found in previous measurements with lighter isoelectronic Si$^{11+}$, Cl$^{14+}$ and also Ti$^{19+}$ ions, [T. Bartsch et al, Phys.\\ Rev.\\ Lett.\\ {\\bf 79}, 2233 (1997); {\\bf 82}, 3779 (1999) and to be published] and in contrast to the results for lighter ions no clear dependence of the electric field enhancement on the magnetic field strength is found. The Maxwellian plasma rate coefficients for $\\Delta n$=0 DR of Ni$^{25+}$ are enhanced by at most $11\\%$ in the presence of the strongest experimentally applied fields. ",
        "introduction": "Dielectronic recombination (DR) is an electron-ion collision process which is well known to be important in astrophysical and fusion plasmas\\cite{DV80,Mue95}. In DR the initially free electron is transferred to a bound state of the ion via a doubly excited intermediate state which is formed by an excitation of the core and a simultaneous attachment of the incident electron. This two step process \\begin{equation} e^- + A^{q+} \\to [A^{(q-1)+}]^{**} \\to A^{(q-1)+}+ h\\nu \\label{eq:DR} \\end{equation} involves dielectronic capture (time-inverse Auger process) as the first step with a subsequent stabilization of the lowered charge state by radiative decay to a state below the ionization limit. This second step competes with autoionization which would transfer the ion back into its initial charge state $q$ with the net effect being resonant elastic or inelastic electron scattering. Another recombination process, which in contrast to DR is non-resonant, is radiative recombination (RR) \\begin{equation} e^- + A^{q+} \\to A^{(q-1)+}+ h\\nu \\label{eq:RR} \\end{equation} where the initially free electron is transferred to a bound state of the ion and a photon is emitted simultaneously. The cross section for RR diverges at zero electron energy and decreases rapidly towards higher energies. In the present investigation we regard RR as a continuous background on top of which DR resonances are emerging.  In the case of narrow non-overlapping DR resonances the DR cross section due to an intermediate state labelled $d$ can be well approximated by\\cite{sho69} \\begin{equation} \\sigma_d(E_{\\rm cm}) = \\bar{\\sigma}_d L_d(E_{\\rm cm}) \\end{equation} with the electron-ion center-of-mass (c.~m.) frame energy $E_{\\rm cm}$, the Lorentzian line shape $L_d(E)$ normalized to $\\int L_d(E) dE=1$ and the resonance strength \\begin{eqnarray} \\bar{\\sigma}_d &=&  4.95\\times10^{-30}{\\rm cm}^2{\\rm eV}^2{\\rm s}\\nonumber \\\\ &&\\times\\frac{1}{E_d} \\frac{g_d}{2 g_i} \\frac{A_{\\rm a}(d\\to i) \\sum_f A_{\\rm r}(d\\to f)} {\\sum_k A_{\\rm a}(d\\to k) + \\sum_{f'}A_{\\rm r}(d\\to f')} \\label{eq:sigma} \\end{eqnarray} where $E_d$ is the resonance energy, $g_i$ and $g_d$ are the statistical weights of the initial ionic core $i$ and the doubly excited intermediate state $d$, $A_{\\rm a}(d\\to i)$ and $A_{\\rm r}(d \\to f)$ denote the rate for an autoionizing transition from $d$ to $i$ and the rate for a radiative transition from $d$ to states $f$ below the first ionization limit, respectively. The summation indices $k$ and $f'$ run over all states which from $d$ can either be reached by autoionization or by radiative transitions, respectively.  Soon after the establishment of DR as an important process governing the charge state balance of ions in the solar corona\\cite{bur64}, Burgess and Summers\\cite{bur69} and Jacobs et al.\\cite{jac76} realized that DR cross sections should be sensitive to external electric fields present in virtually any plasma environment. The electric field enhancement of DR rates was subsequently reproduced in a number of theoretical calculations\\cite{Hahn97}. Briefly, the effect arises from the Stark mixing of $\\ell$ states and the resulting influence on the autoionization rates which, by detailed balancing, determine the capture of the free electron. Autoionization rates strongly decrease with increasing $\\ell$ and, therefore, only low $\\ell$ states significantly contribute to DR. Electric fields mix low and high $\\ell$ states and thereby increase the autoionization rates of the high $\\ell$ states and consequently also the contribution of high Rydberg states to DR.  The control of external fields in experiments using intense ion and electron beams is a challenge. Results from early recombination experiments\\cite{DS92,And92} could only be brought in agreement with theory under the assumption that external electric fields had been present in the interaction region. The first experiment where external fields were applied under well controlled conditions was performed by M\\\"{u}ller and coworkers\\cite{Mueller86/87} who investigated DR in the presence of external fields (DRF) of singly charged Mg$^+$ ions. They observed an increase of the measured DR cross section by a factor of about 1.5 when increasing the motional $\\vec{v}\\times\\vec{B}$ electric field from 7.2 to 23.5 V/cm. The agreement of these results with theoretical predictions\\cite{LNH86,BGP86} was at the 20\\% level. Further DRF experiments with multiply charged C$^{3+}$ ions also revealed drastic DR rate enhancements by electric fields\\cite{Young94}; however, the large uncertainties of these measurements left ambiguities.  The first DRF experiment using highly charged ions at a storage ring was carried out with Si$^{11+}$ ions by Bartsch et al.\\cite{bar97}. It produced results with an unprecedented accuracy, enabling a detailed comparison with theory. Whereas the overall agreement between experiment and theory for the magnitude of the effect (up to a factor of 3 when increasing the field from 0 V/cm to 183 V/cm) was fair, discrepancies remained in the functional dependence of the rate enhancement on the electric field strength. This finding stimulated theoretical investigations of the role of the additional magnetic field which is always present in storage ring DR experiments, since it is needed to guide and confine the electron beam within the electron cooler. In a model calculation Robicheaux and Pindzola\\cite{rob97} found that in a configuration of crossed $\\vec{E}$ and $\\vec{B}$ fields indeed the magnetic field influences  through the mixing of $m$ levels the rate enhancement generated by the electric field. More detailed calculations\\cite{gri98,rob98} confirmed these results. It should be noted that in theoretical calculations by Huber and Bottcher\\cite{HB80} no influence of a pure magnetic field of at least up to 5 T on DR was found.  Inspired by these predictions we previously performed storage ring DRF experiments using Li-like Cl$^{14+}$\\cite{BSM99} and Ti$^{19+}$\\cite{bar99} ions and crossed $\\vec{E}$ and $\\vec{B}$ fields where we clearly discovered a distinct effect of the magnetic field strength on the magnitude of the DR rate enhancement. The electric field effect decreased monotonically with the $\\vec{B}$ field increasing from 30~mT to 80~mT. A decrease of the electric field enhancement by a crossed magnetic field is also predicted by the model calculation of Robicheaux and Pindzola\\cite{rob97} for magnetic fields larger than approximately 20 mT where, due to a dominance of the magnetic over the electric interaction energy, the $\\ell$-mixing weakens and consequently the number of states participating in DR decreases. At lower magnetic fields $m$-mixing yields an increase of the DR rate with increasing $\\vec{B}$ field. A corresponding experimental observation has been made recently by Klimenko and coworkers\\cite{KKG99} who studied recombination of Ba$^+$ ions from a continuum of finite bandwidth which they had prepared by laser excitation of neutral Ba atoms. For a given electric field strength of 0.5 V/cm, they find that the recombination rate is increasingly enhanced by crossed magnetic fields up to about 20 mT. However, there is no effect of the magnetic field when it is directed parallel to the electric field vector. For the $m$-mixing to occur the crossed $\\vec{E}$ and $\\vec{B}$ arrangement is essential. In the case of parallel $\\vec{B}$ and $\\vec{E}$ fields $m$ remains a good quantum number and no influence of the magnetic field is expected.  The aim of the present investigation with Li-like Ni$^{25+}$ is to extend the previous studies to an ion with even higher nuclear charge $Z$. Because of the $Z^4$ scaling of radiative rates it is expected that with higher $Z$ less $\\ell$ states of a given Rydberg $n$ level take part in DR and therefore the sensitivity to $\\ell$-mixing decreases\\cite{Mueller86/87}. Results of Griffin and Pindzola\\cite{gri87} who calculated decreasing DR rate enhancements for increasing charge states of iron ions point into the same direction. Another aspect of going to higher $Z$ is that even high lying $2p_j n\\ell$ Rydberg resonances are more separated in energy and therefore easier to resolve. This has been demonstrated by Brandau et al.\\ who resolved $2p_{1/2}n\\ell$ DR resonances up to $n=41$ in the recombination spectrum of Li-like Au$^{76+}$ ions\\cite{BBD98}. When choosing Ni$^{25+}$ we hoped to be able to study the field enhancement effect on a single $2p_jn\\ell$ resonance. This would enable a quantitative comparison with theory which at present is limited to a single low value of $n$\\cite{gri98} when explicitly treating all $n\\ell m$ levels required for a realistic description of DR in crossed $\\vec{E}$ and $\\vec{B}$ fields.  In our DRF studies we have chosen Li-like ions as test systems since on the one hand their electronic structure is simple enough to be treated theoretically on a high level of sophistication and on the other hand provides strong DR channels connected to the $2s\\to2p_j$ core excitations. ",
        "conclusions": "The recombination of lithium-like Ni$^{25+}$ ions has been experimentally studied in detail. Spectroscopic information on individual DR resonances associated with Ni$^{24+}(1s^22p_jn\\ell)$ intermediate states, which have been experimentally resolved up to $n=32$, has been extracted and the $\\Delta n=0$ DR plasma rate coefficient has been derived. Our experimental result is underestimated by up to a factor of 2 by semi-empirically calculated rate coefficients. At plasma temperatures above 10 eV, results of detailed theoretical calculations are available\\cite{rom88} which agree well with the experiment. At lower temperatures, where Ni$^{25+}$ ions may exist in photoionized plasmas, no data for the DR rate coefficient have been previously available.  In the presence of external electric fields up to 300~V/cm, the measured DR resonance strength is enhanced by a factor 1.5; a rather weak effect in comparison with previous measurements at $Z = 11$, 17 and 22. Due to the overall weakness of the field effect for Ni$^{25+}$ ions, a marked dependence of the DR rate enhancement on the strength of a crossed magnetic field as observed for Cl$^{14+}$ and Ti$^{19+}$ ions has not been detectable in the present investigation.  Experimental limitations prevented us from obtaining an energy resolution in our DRF measurements comparable to that achieved in the field-free measurement. Consequently, we could not resolve a Ni$^{24+}(1s^22p_jn\\ell)$ DR resonance with $n$ high enough to exhibit a field effect. Such an observation would have facilitated a direct comparison with {\\em ab initio} calculations, which due to the large number of $n\\ell m$ states to be considered, presently can only treat a single $n$ manifold of Rydberg states in the presence of crossed $\\vec{E}$ and $\\vec{B}$ fields\\cite{gri98}. Improvements of this situation can be expected in the near future from the steady increase of computing power on the theoretical side, and on the experimental side from a dedicated electron target which is presently being installed at the TSR. With the electron target and the electron cooler operating at the same time we will be able to perform DRF measurements with continuously cooled ion beams, yielding DRF spectra with increased resolution."
    },
    "0002/astro-ph0002137_arXiv.txt": {
        "abstract": "Blazars are known to emit a broad band emission from radio to gamma-rays with rapid time variations, particularly, in X- and gamma-rays.  Synchrotron radiation and inverse Compton scattering are thought to play an important role in emission and the time variations are likely related to the acceleration of nonthermal electrons. As simultaneous multiwavelength observations with continuous time spans are recently available, some characteristics of electron acceleration are possibly inferred from the spectral changes of high energy emission. In order to make such inferences, we solve the time-dependent kinetic equations of electrons and photons simultaneously using a simple model for electron acceleration. We then show how the time variations of emission are dependent on electron acceleration. We also present a simple model for a flare in X-rays and TeV gamma-rays by temporarily changing the acceleration timescale. Our model will be used, in future, to analyze observed data in detail to obtain information on electron acceleration in blazars. ",
        "introduction": "High energy emission from blazars is usually thought to be produced by relativistically moving jets or blobs from the nucleus of galaxies \\citep[e.g.,][]{br78,bk79,mgc92,sbr94,it96}. The physical properties of such jets have been probed mostly based on the steady state models of synchrotron radiation and inverse Compton scattering by a nonthermal electron population (SSC model). However, blazars are also characterized by rapid and strong time variability. Recent observations have revealed that the emission exhibits short time variations in X- and gamma-ray bands on timescales from weeks down to half an hour \\citep[e.g.,][for review]{muk97,umu97}, as fast time variations of Mrk 421 were observed by X-rays and TeV gamma-rays \\citep{gaid96,taka96}; similar time variations of Mrk 501 were also found by multiwavelength observations \\citep{kat99}. These observations should provide important clues on physical processes in relativistic jets, in particular, on electron acceleration.  To make theoretical inferences, we need to calculate time-dependent emission spectra from a time-dependent electron population. An example of such theoretical models was recently presented by \\cite{mk97}.  They solved the kinetic equations of electrons and photons simultaneously, by injecting power-law electrons with an exponential cutoff.  They showed various possibilities to explain the time variations observed from Mrk 421, such as the changes in the magnetic field or the maximum Lorentz factor of nonthermal electrons.  Although the correlation between X-rays and TeV gamma-rays are important to discuss the SSC model, their estimate of Compton scattering in the Klein-Nishina regime does not necessarily correctly account for the energy change in scatterings, because of a simplified treatment of Compton scattering in the Klein-Nishina regime [see \\cite{mk95} for the details of their calculation method]. \\cite{kirketal98}, on the other hand, extended the above model, assuming that an acceleration region and a cooling region are spatially separated; i.e., electrons accelerated in a shock region are transferred to a cooling region where they emit synchrotron photons (they did not include Compton scattering). Their model was intended to explain the time variability of X-rays, by changing acceleration timescale.  Besides the time variations of flare activities explained by \\cite{mk97} and \\cite{kirketal98}, the early stages of acceleration are of great importance. By examining the properties of the time evolutions of photon spectra during acceleration, we may obtain the diagnoses of acceleration mechanisms. The recent development of observations in X-rays (ASCA and Beppo-SAX) to TeV gamma-rays (e.g., Whipple and HEGRA) and future experiments might be used to confirm the diagnoses.  In this paper, we use a formulation similar to \\cite{mk97}, but with the full Klein-Nishina cross section in the Compton scattering kernel, so that the emission in the TeV range is calculated correctly. We also include a particle acceleration process by considering spatially separated acceleration and cooling regions as in \\cite{kirketal98}, although we do not consider the spatial transfer of electrons. We particularly emphasize the detailed study of the properties of electron and photon spectra in the early stage of acceleration.  We describe our model in \\S \\ref{sec:model} and present numerical results in \\S \\ref{sec:results}. Summary of our results is given in \\S \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  Simulations of the time evolution of electron and photon energy distributions were presented as a model of time variations observed by X- and gamma-rays from blazars. By assuming that acceleration and cooling regions in a blob are spatially separated, we calculated the energy spectra of electrons in each regions. Electrons in the acceleration region are accelerated with a characteristic timescale $t_{\\rm acc}$ and escape on a timescale $t_{e, {\\rm esc}}$; here we assumed $t_{\\rm acc} = t_{e, {\\rm esc}}$, so that the electron spectrum in a steady state obeys a power law, $N(\\gamma) \\propto \\gamma^{-2}$, as realized in the standard model of shock acceleration \\citep[e.g.,][]{druly,be87}. Electrons escaping from the acceleration region are injected into the cooling region where they lose energy by radiation and finally escape from the blob on a timescale assumed to be $2R/c$. With these assumptions, we performed the simulations of the time evolutions of electrons and photons for various values of parameters. Although we did not include a specific acceleration mechanism, we took into account the salient features of diffusive shock acceleration, so that we could study the properties of time variation accompanying shock acceleration.  We first presented the results of the time evolution of the spectral energy distribution of radiation associated with the evolution of the electron number spectrum. In the early stage of the evolution, i.e., $t = 0$ -- $R/c$, the synchrotron component dominates the spectrum. The energy flux of soft X-rays starts to rise earlier than that of hard X-rays. Later ($t > R/c$), the Compton luminosity gradually increases. At the same time, the peak energy of the synchrotron component decreases because of radiative cooling. It was found that in a steady state, escaping electrons carry more energy than radiation: This result, of course, depends on the values of the parameters used. We also showed the dependence of time evolution on the acceleration timescale, the electron injection rate, and the strength of magnetic fields. The value of $\\gamma_{\\rm max}$ and the ratio of the synchrotron luminosity to the Compton luminosity depend on such parameters.  We next simulated a flare by simply changing the value of $t_{\\rm acc}$ for a certain time span. With a shorter acceleration timescale, more energetic electrons are produced and consequently more hard photons are produced. The relation between the energy flux and the photon index during a flare was obtained, which is similar to the one observed from Mrk 421 \\citep{taka96}.  Our formulation provides a method to treat high energy flares including particle acceleration processes, which is beyond usual analyses where nonthermal electron spectra are arbitrarily assumed and only cooling processes are included. Although we have not applied our model to any specific case of flares, it is straightforward to do this using our code. The examples presented here seem to cover a wide range of observed flares. These applications are deferred to future work. On the theoretical side, as proposed by \\cite{kirketal98}, electrons accelerated at a shock are transferred outside of the shock and cool radiatively. To include such spatial transfer of electrons, we, in future, need to solve for the structure around acceleration regions.  Recently \\cite{cg99} showed observational consequences associated with time variations with timescales shorter than $R/c$. When such short timescale variations occur, observed emission is a superposition from various parts of a cloud. Then the time profile of each time variation is not necessarily observed clearly. The model presented in this paper contains the acceleration timescale shorter than $R/c$. Thus our model may not directly reflect observed spectra. However, to understand the relation between electron acceleration and time variation of emission, such a study should be useful."
    },
    "0002/astro-ph0002247_arXiv.txt": {
        "abstract": "The thermal evolution of the intergalactic medium (IGM) depends on the reionization history of the universe. Numerical simulations indicate that the low density IGM, which is responsible for the low column density \\lya\\ forest, follows a well defined temperature-density relation. This results in a cut-off in the distribution of line widths as a function of column density. We use hydrodynamic simulations to calibrate the relation between the cut-off and the temperature-density relation and apply this relation to Keck spectra spanning a redshift range $z=2$--4.5. We find that the temperature peaks at $z\\sim 3$ and interpret this as evidence for reheating due to the reionization of helium. ",
        "introduction": "Quasars have provided us with a unique probe of the high redshift universe. These bright point sources shine like a flashlight through space, revealing the presence of baryonic matter through the light it absorbs. Thus every quasar spectrum contains a one-dimensional map of the distribution of matter along the line of sight. The extraordinary quality of the spectra obtainable with the HIRES spectrograph on the Keck telescope, enables us to extract the wealth of information that has been collected by the quasar's light along its journey through space and time. Computer simulations of structure formation have been remarkably successful in reproducing these observations. They show that the physics governing the high redshift intergalactic medium (IGM), which is responsible for the low column density absorption lines (the so-called \\lya\\ forest) is relatively simple. The IGM, which contains most of the baryons in the universe, is photoionized and photoheated by the collective UV radiation from young stars and quasars. On large scales its dynamics are determined by the gravitational field of the dark matter, while on small scales gas pressure is important. The availability of superb data and a detailed physical model, have made the \\lya\\ forest into a powerful probe of the high-redshift universe.  Since shock heating is unimportant in the low-density IGM, most of the gas follows a simple temperature-density relation which is the result of the interplay of photoionization heating and adiabatic cooling due to the expansion of the universe. For densities around the cosmic mean, this relation is well-described by a power-law, $T=T_0(\\rho/\\bar{\\rho})^{\\gamma-1}$ (Hui \\& Gnedin 1997). At reionization the gas is reheated, resulting in an increase in $T_0$ and a decrease in $\\gamma$ (provided that the gas is reionized on a timescale short compared to the Hubble time). In ionization equilibrium, $T_0$ decreases and the slope of the effective equation of state steepens (i.e.\\ $\\gamma$ increases). However, because the timescale for recombination is long, the gas retains some memory of how and when it was reionized (Miralda-Escud\\'e \\& Rees 1994).  The distribution of line widths ($b$-parameters) depends on various mechanisms. Thermal motions of the hydrogen atoms broaden the \\HI\\ absorption lines and other processes, such as the differential Hubble flow across the absorbing structure and bulk flows, also contribute to the line widths. However, the minimum line width is set by the temperature of the gas, which in turn depends on the density. A standard way of analyzing \\lya\\ forest spectra is by decomposing them into a set of Voigt profiles. Since the minimum line width ($b$-parameter) depends on the temperature, and since column density ($N$) correlates strongly with physical density, there is a cut-off in the $b(N)$ distribution which traces the effective equation of state of the IGM (Schaye et al.~1999; Ricotti, Gnedin \\& Shull 2000; Bryan \\& Machacek 2000). We have used this relation to measure the thermal evolution of the IGM from a set of nine \\lya\\ quasar absorption line spectra.  This work is more fully described and discussed in a forthcoming publication (Schaye et al.\\ 2000). ",
        "conclusions": "The measured evolution of the temperature at the mean density and the slope of the effective equation of state are plotted in Figure~1. From $z\\sim 3.5$ to $z\\sim 3.0$, $T_0$ increases and the gas becomes close to isothermal ($\\gamma \\sim 1.0$). This behavior differs drastically from that predicted by models in which helium is fully reionized at higher redshift. For example, the solid lines correspond to a simulation that uses a uniform metagalactic UV-background from quasars as computed by Haardt \\& Madau (1996) and which assumes the gas to be optically thin. In this simulation, both hydrogen and helium are fully reionized by $z\\sim 4.5$ and the temperature of the IGM declines slowly as the universe expands. Such a model can clearly not account for the peak in the temperature at $z\\sim 3$ (reduced $\\chi^2$ for the solid curves are 6.9 for $T_0$ and 3.6 for $\\gamma$). Instead, we associate the peak in $T_0$ and the low value of $\\gamma$ with reheating due to the second reionization of helium (\\HeII\\ $\\rightarrow$ \\HeIII). This interpretation is supported by measurements of the \\SiIV/\\CIV\\ ratio (Songaila 1998, but see also Boksenberg, Rauch, \\& Sargent 1998 and Giroux \\& Shull 1997) and direct measurements of the \\HeII\\ opacity (Heap et al.\\ 2000 and references therein).  \\begin{figure} \\mbox{ \\hspace{-0.8cm} \\plotone{schayej_fig1a.eps} \\hspace{-0.8cm} \\plotone{schayej_fig1b.eps} } \\caption{The evolution of the temperature at the mean density (left panel) and the slope of the effective equation of state (right panel). Horizontal error bars indicate redshift interval spanned by the absorption lines, vertical error bars are $1\\,\\sigma$ errors. Different symbols correspond to different quasars. See text for a description of the models.} \\end{figure}  The dashed lines in Figure~1 are for a model that was designed to fit the data (reduced $\\chi^2$ is 0.24 for $T_0$ and 1.38 for $\\gamma$). In this simulation, which has a much softer UV-background at high redshift, \\HeII\\ reionizes at $z\\sim 3.2$. Before reionization, when the gas is optically thick to ionizing photons, the mean energy per photoionization is much higher than in the optically thin limit (Abel \\& Haehnelt 1999). We have approximated this effect in this simulation by enhancing the photoheating rates during reionization, so raising the temperature of the IGM.  Since the simulation assumes a uniform ionizing background, the temperature has to increase abruptly (i.e.\\ much faster than the gas can recombine) in order to make $\\gamma$ as small as observed. In reality, the low-density gas may be reionized by harder photons, which will be the first ionizing photons to escape from the dense regions surrounding the sources. This would lead to a larger temperature increase in the more dilute, cooler regions, resulting in a decrease of $\\gamma$ even for a more gradual reionization. Furthermore, although reionization may proceed fast locally (as in our small simulation box), it may be patchy and take some time to complete. Hence the steep temperature jump indicated by the dashed line, although compatible with the data, should be regarded as illustrative only. The globally averaged $T_0$ could well increase more gradually which would also be consistent with the data. More data at $z \\ga 3$ is needed to determine whether the temperature rise is sharp or gradual. On the theoretical side, more realistic models should include radiative transfer effects, which are important during reionization.  Together with measurements of the \\HeII\\ opacity, which probe the ionization state in the voids, the thermal history of the IGM provides important constraints on models of helium reionization. Furthermore, the temperature of the IGM before the onset of helium reionization can be used to constrain the redshift of hydrogen reionization, which marks the end of the dark ages of cosmic history."
    },
    "0002/astro-ph0002071_arXiv.txt": {
        "abstract": "Standard big bang nucleosynthesis (SBBN) has been remarkably successful, and it may well be the correct and sufficient account of what happened.  However, interest in variations from the standard picture come from two sources: First, big bang nucleosynthesis can be used to constrain physics of the early universe.  Second, there may be some discrepancy between predictions of SBBN and observations of abundances. Various alternatives to SBBN include inhomogeneous nucleosynthesis, nucleosynthesis with antimatter, and nonstandard neutrino physics. ",
        "introduction": "The success of standard big bang nucleosynthesis (SBBN) in predicting the observed abundances of the light elements has led to the widespread view that SBBN must be correct.  According to this view, any remaining disagreements must be due to systematic errors in observations or incorrect, or too crude, chemical evolution models.  While this view may well be the right one, we should not be blind to other possibilities.  I will stay within the context of the Hot Big Bang (for an alternative, see Burbidge \\& Hoyle 1998), and discuss some models of nonstandard big bang nucleosynthesis (NSBBN).  NSBBN scenarios range from small modifications to SBBN to a complete change in the decisive physical phenomena, like in the late-decaying massive particle scheme of Dimopoulos et al. (1988).  Motivations for studying NSBBN go in two directions.  First, the remarkable success of SBBN allows one to severely constrain the physics of the early universe.  If one tries to change the conditions from the standard assumptions the resulting abundances of the light elements differ from the observed ones.  For many modifications, BBN provides the strongest constraints.  BBN gives also the strongest constraint on the single parameter of SBBN, the baryon density, usually given as the baryon-to-photon ratio, \\begin{equation} \\eta \\equiv \\frac{n_b}{n_\\gamma}, \\qquad \\eta_{10}\\equiv10^{10}\\eta. \\end{equation} Second, one may try to improve on SBBN.  From time to time it has seemed that there might be some discrepancy between observations and SBBN, which could then be explained by NSBBN. In particular, there has been tension between D/H and $Y_p$ (see, e.g., Hata et al. 1995).  To relieve this tension, either a lower D or a lower ${}^4$He yield has been looked for. Also one may want to relax the SBBN bounds to $\\eta$. Other astronomical considerations have given motivation for trying to raise the upper limit to $\\eta$. If one believes that the energy density of the universe is dominated by vacuum energy (the cosmological constant) and accepts the newer observations on D/H and $Y_p$ favoring somewhat larger $\\eta$ within SBBN, this motivation largely disappears.  There is a very large body of work on NSBBN. Extensive reviews are given by Malaney \\& Mathews (1993) and Sarkar (1996), which contain, respectively, over 500 and over 700 references.  Here I will be able to mention only a random few.  Most of the work on NSBBN can be divided into four broad classes: \\begin{enumerate} \\item Inhomogeneous BBN.  Usually this means inhomogeneity in the baryon-to-photon ratio, $\\eta$, but there are also other possibilities, like inhomogeneity in the neutrino chemical potentials. \\item Nonstandard neutrino physics, e.g., additional (``sterile'') neutrino flavors, neutrino degeneracy (asymmetry), massive $\\nu_{\\tau}$, or neutrino oscillations. \\item Late-decaying ($\\tau$ = 1--$10^8$ s) massive particles, black holes, cosmic strings,  etc. \\item Time-varying fundamental constants. \\end{enumerate} In the interest of time and space, I will discuss the first two classes only. ",
        "conclusions": "At present, no NSBBN scenario appears as convincing as SBBN, which is the simplest of all. Often the real world has turned out to be more complicated in the end than first assumed, but for the early universe a simple picture has been very successful.  However, it is healthy to keep in mind the possibility that SBBN might not be the full story, and that any discrepancies between observations and SBBN might actually be telling us something important about the early universe or particle physics."
    },
    "0002/gr-qc0002036_arXiv.txt": {
        "abstract": "\\hfill{\\small\\bf Abstract}\\hfill\\smallskip \\par We investigate the propagation of electromagnetic waves through a static wormhole. It is shown that the problem can be reduced to a one-dimensional Schr\\\"odinger-like equation with a barrier-type potential. Using numerical methods, we calculate the transmission coefficient as a function of the energy. We also discuss the polarization of the outgoing radiation due to this gravitational scattering. ",
        "introduction": "Since the pioneering article by Morris and Thorne \\cite{motho}, Lorentzian wormholes have attracted a lot of interest in the literature. Let us recall that a wormhole is a solution to Einstein's equations that can be understood as a ``handle'' connecting either two universes or two distant places in the same universe.  A number of static and dynamic wormholes have been found both in General Relativity and in alternative theories of gravitation (see \\cite{hochvis1,hochvis2} and references therein).  It has also been shown that under particular circumstances this geometry permits the formation of closed timelike curves (CTC's) \\cite{timemach}.  One of the most striking features of this field configuration is that it needs matter with negative energy density as a source \\cite{motho}.  Hochberg and Visser, and Ida and Hayward have analyzed in detail the issue of the violation of the energy conditions both in static \\wh s \\cite{hochvis1} and in a completely general (\\ie\\/  non-symmetric and time dependent) traversable \\wh\\/ \\cite{hochvis2,ida}.  Their analysis shows that the null energy condition (NEC) must be violated in both cases.  Since it is believed that all classical matter acting as a source of gravitation satisfies the NEC, only quantum phenomena can, in principle, be responsible for NEC violations.  In fact, it has been known for some time that there are states in Quantum Field Theory that may violate the energy conditions \\cite{epstein}.  Many examples of systems that can violate these conditions due to quantum effects can be found in \\cite{visserlib}; one such example is the conformally coupled scalar field. Indeed, Hochberg \\et\\/ \\cite{hochberg} found the first self-consistent \\wh\\/ solution in semiclassical General Relativity with quantized conformally coupled scalar matter as a source \\cite{bar}. Another instance in which quantum effects are essential has been presented recently by Kim and Lee \\cite{kimlee}. Using a two-dimensional dilaton-gravity at the one-loop level, they showed that a wormhole can be the final state of an evaporating black hole. This result agrees with previous considerations of Hayward \\cite{sean}, and might be considered as an example of a process to bring a wormhole into existence. Recently, Krasnikov \\cite{kras} showed that there exist static wormholes with the vacuum stress-energy tensor of the neutrino, the electromagnetic or the massless sclar field as a source of Einstein's equations.  In spite of these results, the issue of the existence of large amounts of NEC violating matter has not yet been resolved.  However, one can take a complementary view by assuming that wormholes exist and work out possible consequences.  This line of reasoning was started by Frolov and Novikov \\cite{frono}, who studied nontrivial effects arising in a \\wh\\/ that interacts with an external electromagnetic or gravitational field.  In a subsequent paper \\cite{frono2}, they used wormholes as tools for studying the interior of black holes.  Later, Gonzalez-Diaz studied some astrophysical consequences of the existence of these objects \\cite{gondiaz}.  More recently, Cramer \\et \\cite{cramer} considered the unusual features of the lensing of light caused by an object with negative mass (which could be interpreted as one of the mouths of a \\wh ).  As a natural application of this idea, Torres \\et\\/ \\cite{torres} considered the microlensing caused by wormholes on light coming from  distant active galactic nuclei.  They showed that these events resemble certain types of Gamma Ray Bursts (GRB's) and set an upper bound on the negative mass density existing in the universe in the form of wormholes. The GRB's produced by microlensing by wormholes present a definite feature that differentiate them from those associated with fireballs \\cite{fireball}. Namely, they always appear in pairs called FRED-antiFRED. A subsequent article by Anchordoqui {\\em et al} \\cite{doqui2} investigated profiles of GRB's from the BATSE database, searching for observable signatures of natural wormholes.  They could identify at least one event that may be associated with this mechanism.  From an observational point of view, it is also important to study the propagation of different types of perturbations in the geometry associated with these objects.  This study was initiated by Kar and Sahdev \\cite{kar1}, and Kar \\et\\/ \\cite{kar2}, who studied the reflection and transmission of massless scalar waves in the presence of an ultrastatic wormhole.  In this article we consider the propagation of electromagnetic waves in the same geometry.  We will study some properties of the outgoing radiation modified by the gravitational field of the wormhole.  We should mention at this point an article that is at first sight related to ours. In \\cite{clement}, Cl\\'ement discussed the problem of the scattering of scalar and electromagnetic waves in an Ellis geometry (which is in fact a wormhole). Following the ideas of Wheeler's geometrodynamics \\cite{wheeler}, his aim was to show that wormholes are particle-like objects. Consequently, the wormhole structure in \\cite{clement} was to be important only at a microscopic level, while in this work we are considering macroscopic wormholes.  The structure of the paper is as follows.  In the next section we transform Maxwell's equations to a form that is convenient to study the propagation of electromagnetic waves in curved spacetimes.  Following this, we introduce the metric of a static wormhole in which we will study the electromagnetic perturbations.  Section IV shows how the problem is equivalent to a one-dimensional problem in the presence of a potential barrier.  We also present numerical results obtained for the transmission coefficient.  Section V investigates the polarization of the outgoing radiation.  We conclude with a summary of our results. ",
        "conclusions": "We have studied here several aspects of the transit of electromagnetic waves through an ultrastatic wormhole. We have shown that the problem can be rewritten in such a way that the curved background geometry is replaced by a medium whose properties are summarized by the refraction index $n(\\rho)$. We also showed that the magnetic modes of the electromagnetic field obey a one-dimensional Schr\\\"odinger-like equation, with an asymmetric barrier-type potential. The transmission coefficient was calculated numerically, and it exhibits some features common to barrier-like potentials.  Namely, in a small interval of $k$, the curve rapidly increases until there is no reflection and for higher values of $J$, initially the curve rises more slowly.  We have also demonstrated that the interaction of the radiation with the gravitational field of the wormhole cannot change the polarization state of the radiation. Moreover, this was shown for any static wormhole and not only with the particular case described by metric Eq.(\\ref{met2}).  It would be interesting to search for special features of the transmission coefficient in more general types of wormholes. In particular, the existence of bound states (\\ie standing waves) for scalar waves was shown in \\cite{kar2}. If  resonances with the same characteristics also exist in the case of electromagnetic waves, their position in what could be considered as the emission spectrum of the wormhole may give us information on the size of the throat and its shape \\cite{kar2}.  One can expect a greater richness of observational consequences if the background wormhole geometry is rotating \\cite{theo}. For instance, differential gravitational scattering of polarized radiation can be caused by the  helicity-rotation coupling \\cite{mash3}. Finally, to complete the study of perturbations in these spacetimes, it would be necessary to analyze spinor and gravitational perturbations. We intend to address these problems in future works."
    },
    "0002/astro-ph0002357_arXiv.txt": {
        "abstract": "We present infrared spectroscopy of the Antennae Galaxies (NGC~4038/4039) with NIRSPEC at the W. M. Keck Observatory.  We imaged the star clusters in the vicinity of the southern nucleus (NGC~4039) in $0\\arcs.39$ seeing in K-band using NIRSPEC's slit-viewing camera. The brightest star cluster revealed in the near-IR (M$_{\\rm K}(0) \\simeq -17.9$) is insignificant optically, but coincident with the highest surface brightness peak in the mid-IR ($12-18 \\mu$m) ISO image presented by \\citet{mirabel98}.  We obtained high signal-to-noise 2.03$-$2.45 $\\mu$m spectra of the nucleus and the obscured star cluster at R $\\sim 1900$.  The cluster is very young (age $\\sim 4$ Myr), massive (M $\\sim 16 \\times 10^6$ M$_{\\odot}$), and compact (density $\\sim 115$ M$_{\\odot}$ pc$^{-3}$ within a 32 pc half-light radius), assuming a Salpeter IMF (0.1$-$100 M$_{\\odot}$).  Its hot stars have a radiation field characterized by T$_{\\rm eff}\\sim 39,000$ K, and they ionize a compact \\ion{H}{2} region with n$_{\\rm e}\\sim 10^4$ cm$^{-3}$.  The stars are deeply embedded in gas and dust (A$_{\\rm V} \\sim 9-10$ mag), and their strong FUV field powers a clumpy photodissociation region with densities n$_{\\rm H}\\ga 10^5$ cm$^{-3}$ on scales of $\\sim 200$ pc, radiating L$_{{\\rm H}_2 1-0~{\\rm S}(1)}= 9600$ L$_{\\odot}$. ",
        "introduction": "The Antennae (NGC~4038/4039) are a pair of disk galaxies in an early stage of merging which contain numerous massive super star clusters (SSCs) along their spiral arms and around their interaction region \\citep{whitmore95,whitmore99}.  The molecular gas distribution peaks at both nuclei and in the overlap region \\citep{stanford90}, but the gas is not yet undergoing a global starburst typical of more advanced mergers \\citep{nikola98}.  Star formation in starbursts appears to occur preferentially in SSCs.  We chose to observe the Antennae because their proximity permits an unusually detailed view of the first generation of merger-induced SSCs and their influence on the surrounding interstellar medium.  The Infrared Space Observatory (ISO) $12-18~\\mu$m image showed that the hot dust distribution is similar to that of the gas, but peaks at an otherwise inconspicuous point on the southern edge of the overlap region \\citep{mirabel98}.  This powerful starburst knot is also a flat-spectrum radio continuum source \\citep{hummel86} and may be associated with an X-ray source \\citep{fabbiano97}.  We imaged the region around this knot, and discovered a bright compact star cluster coincident with the mid-IR peak.  We obtained moderate-resolution (R $\\sim 1900$) K-band spectra of both the obscured cluster and the NGC~4039 nucleus. ",
        "conclusions": "The highest surface brightness mid-IR peak in the ISO map of the Antennae Galaxies is a massive ($\\sim 16\\times$10$^6$ M$_{\\odot}$), obscured (A$_{\\rm V} \\sim 9-10$), young (age $\\sim 4$ Myr) star cluster with half-light radius $\\sim$ 32 pc, whose strong FUV flux excites the surrounding molecular ISM on scales of up to 200 pc.  The cluster spectrum is dominated by extended fluorescently excited H$_{2}$ emission from clumpy PDRs and nebular emission from compact \\ion{H}{2} regions.  In contrast, the nearby nucleus of NGC~4039 has a strong stellar spectrum dominated by cool stars, where the only emission lines are due to shock-excited H$_2$.  These observations confirm the potential of near-infrared spectroscopy for exploration and discovery with the new generation of large ground-based telescopes.  Our ongoing program of NIRSPEC observations promises to reveal a wealth of information on the nature of star formation in star clusters.  \\vspace{-0.05in} {\\plotfiddle{chi2temp_col.eps}{2.55in}{0}{100}{100}{0}{10}} {\\footnotesize Figure 4. - Comparison of H$_2$ line strengths with PDR models. Contours of $\\chi^2_\\nu$ for 22 lines projected onto n$_{\\rm H} -$G$_0$ plane peak at n$_{\\rm H} \\sim 10^5$ cm$^{-3}$ and G$_0 \\sim 5000$. Model points (+) are for T$_0$ = 300 $-$ 2000 K. White + marks best-fit PDR model of Draine \\& Bertoldi, with T$_0$ = 900 K and $\\chi_\\nu^2=9.3$.  Contours are 50, 25, 20, 15, 12, 10. }"
    },
    "0002/astro-ph0002161_arXiv.txt": {
        "abstract": "We present preliminary results from a Hubble Space Telescope (HST) WFPC2 investigation of spatial and temporal distributions of star clusters in the clumpy irregular galaxy NGC~7673 and the starburst spirals NGC 3310 and Haro~1.   We compare the spectral energy distributions of star clusters in the large clumps in NGC 7673 to model calculations of stellar clusters of various ages. We also propose that the presence of super star clusters in clumps seems to be a feature of intense starbursts. ",
        "introduction": "The optical morphologies of active star forming galaxies are often dominated by kpc-scale star-forming regions, which appear as luminous `clumps' or `islands'. These clumps can have very high optical and UV surface brightnesses, making them visible signatures of intense star formation over cosmological distances.  High angular resolution imaging with the {\\it Hubble Space Telescope} (HST) and ground-based telescopes, reveals that clumps consist of associations of dense star clusters. These are usually superimposed on a more diffuse background, probably made up of massive stars.  The presence of numerous  star clusters influences the evolution of clumps through their interactions with the surrounding ISM, that must respond to photoionization, as well as mechanical energy and momentum inputs from the evolving star clusters. These processes can trigger star formation and simultaneously remove ISM from the clump. Since star clusters can be age-dated from their spectra, their presence also offers the possibility of measuring the history of star-forming activity in starburst galaxies.  Starbursts are well-known producers of dense, massive star clusters and thus are an excellence place to study these interactions. In this paper we explore the spatial and temporal patterns of massive star cluster formation in a small sample of relatively nearby starburst galaxies.  The primary objects in this study are NGC 2415 (Haro 1), NGC 3310 and NGC 7673. The Haro 1 and NGC 7673 data are HST WFPC2 ultraviolet (UV) and optical images, while we use ground-based optical images from the WIYN 3.5-m telescope for NGC 3310. Details are presented in Gallagher et al. (2000).  \\begin{table} \\begin{center} Properties of Studied Starbursts \\begin{tabular}{rlll} \\hline Galaxy & Distance & Magnitude & Size   \\\\ \\hline \\hline NGC 2415    & D=50 Mpc &  M$_B$=-20.8 &  A=13 kpc \\\\  NGC 3310    &       D=19 Mpc & M$_B$=-19.9 & A=17 kpc \\\\  NGC 7673    &       D=46 Mpc &  M$_B$=-20  &  A=17 kpc \\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "Clumps seem to be features of active star bursts with high intensities. This fits with theoretical models (Elmegreen \\& Efrenov 1997; Noguchi 1999) where stellar clumps form in gas-rich galactic disks with higher than average internal velocity dispersions, leading to large Jeans masses and lengths. Such conditions could be natural consequences of collisional perturbations of gassy disk galaxies. They would be more common in less evolved galaxies with higher gas contents and lower degrees of organization, providing a possible explanation for the clumpy appearances of high redshift galaxies (Noguchi 1997). Although super star clusters can form in a range of conditions, clumps may be particularly important in unevolved galaxies; observations of nearby clumps then may provide insight into the formation of globular clusters.  Once star formation begins in a clump, the subsequent evolution must be complex. Energy and momentum inputs from massive stars and star clusters produce obvious observational signatures in emission line profiles (Homeier \\& Gallagher 1999) will disturb the ISM, and likely lead to compressed regions which are excellent sites for further cluster formation (eg., Scalo \\& Chappell 1999). The close spacing between clusters and high densities of gas clouds may also lead to cluster-cluster mergers or cluster rejuvenations due to gas cloud captures. These possibilities reinforce the capabilities for massive star forming clumps to be prime producers of massive super star clusters. Once made these clusters will be the most likely to survive and observed at later times.  We thank the WFPC2 Investigation Definition Team, and the Space Telescope Science Institute Archival Observer program for providing support through the National Aeronautics and Space Administration for this work."
    },
    "0002/astro-ph0002482_arXiv.txt": {
        "abstract": "We present two simple dynamical models for Sagittarius based on N-body simulations of the progressive disruption of a satellite galaxy orbiting for 12.5 Gyr within a realistic Galactic potential.  In both models the satellite initially has observable properties similar to those of current outlying dwarfs; in one case it is purely stellar while in the other it is embedded in an extended massive halo.  The purely stellar progenitor is a King model with a total velocity dispersion of 18.9 $\\kms$, a core radius of 0.44 kpc and a tidal radius of 3 kpc.  The initial stellar distribution in the other case follows a King profile with the same core radius, a slightly larger total velocity dispersion and similar extent. Both these models are consistent with all published data on the current Sagittarius system, they match not only the observed properties of the main body of Sagittarius, but also those reported for unbound debris at larger distances. ",
        "introduction": "The Sagittarius dwarf galaxy is the closest satellite of the Milky Way (Ibata, Gilmore \\& Irwin 1994, 1995, hereafter IGI95). Soon after its discovery, several groups carried out simulations to see if its properties are consistent with the disruption of an object similar to the other dwarf companions of the Milky Way, but none produced a model in full agreement with both the age and the structure of the observed system (Johnston, Spergel \\& Hernquist 1995; Vel\\'azquez \\& White 1995; Edelsohn \\& Elmegreen 1997; Ibata et al. 1997, hereafter I97; G\\'omez-Flechoso, Fux \\& Martinet 1999).  All groups assumed light to trace mass and an initial system similar to observed dwarf spheroidals. All found the simulated galaxy to disrupt after one or two orbits whereas the observed system has apparently completed ten or more. Most considered this to be a problem (but cf Vel\\'azquez \\& White 1995). As a result, several unconventional models were proposed to explain the survival and structure of Sagittarius.  In an extensive numerical study, Ibata \\& Lewis (1998) concluded that Sagittarius must have a stiff and extended dark matter halo if it is to survive with 25\\% of its initial mass still bound today.  Since an extended halo cannot remain undistorted in the Galaxy's tidal field for any conventional form of dark matter, it is unclear how this idea should be interpreted. Furthermore, it produces an uncomfortably large mass-to-light ratio ($\\sim 100$), it cannot reproduce the observed elongation, and it suggests that little tidal debris will be liberated, in apparent conflict with the observations of Mateo, Olszewski \\& Morrison (1998), and Majewski et al. (1999) (see also Johnston et al. 1999). A somewhat less unorthodox model was proposed by Zhao (1998), where Sagittarius was scattered onto its current tightly bound orbit by an encounter with the Magellanic Clouds about 2 Gyr ago. This appears physically possible but requires careful tuning of the orbits of the two systems (see Ibata \\& Lewis 1998; and Jiang \\& Binney 2000).  Another mechanism by which the dwarf could have moved to a short-period orbit is dynamical friction, which can be important only if Sagittarius has lost a lot of mass in the past.  Jiang \\& Binney (2000) found a one-parameter family of initial configurations that evolve into something like the present system over a Hubble time. Their initial systems have masses $\\sim10^{10-11}\\sm$ and start from a Galactocentric radius $\\sim200$ kpc.  Driven by this apparent puzzle, we decided to search more thoroughly for a self-consistent model of the disruption of Sagittarius, which, after a Hubble time, has similar characteristics to those observed. (See Table~\\ref{table_sag} for a summary of the observed properties of the system.) Below we present two models which meet these requirements. \\begin{table} \\caption{Properties of Sagittarius \\, (IGI95, I97)} \\label{table_sag} \\begin{tabular}{lc} \\hline Orbital properties & \\\\ \\hline distance from the Sun $d$ & $25 \\pm 2$ kpc \\\\ heliocentric radial velocity $v_r^{\\rm sun}$ & $140 \\pm 2 \\kms$ \\\\ proper motion in $b$ $\\mu_{b}$ & $250 \\pm 90 \\kms$ \\\\ gradient along the orbit ${\\rm d}v_r/{\\rm d}b$ & $< 3 \\kms/\\deg$ \\\\ angular position in the sky $(l, b)$ & $(5.6\\ndeg,-14\\ndeg)$ \\\\ \\hline Internal properties & \\\\ \\hline luminosity & $\\gtrsim 10^7 \\sL$ \\\\ velocity dispersion $\\sigma(v_r)$ & $11.4 \\pm 1 \\kms$ \\\\ angular extent in ($l$, $b$) & $8 \\ndeg \\times 3 \\ndeg$ \\\\ half-mass radius & 0.55 kpc \\\\ mean metallicity $\\langle$[Fe/H]$\\rangle$ & $\\sim -1.$ dex \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "We have found viable models for the Sagittarius dwarf galaxy with a wide range of total luminosities and masses, and both with and without extended dark halos. A purely stellar progenitor could be a King model with a total velocity dispersion of 18.9~$\\kms$, a core radius of 0.44~kpc and a tidal radius of 3~kpc. For the case where the progenitor is embedded in an extended massive halo, the initial stellar distribution follows a King profile with the same core radius, a slightly larger total velocity dispersion of $23 \\kms$ and similar extent.  The dark-matter is more extended. The data available at present only weakly constrain the total initial extent either of the light or of the mass. The observed metallicity data, for example, are consistent with an initial galaxy similar to either of our detailed models, both of which would lie within the scatter of the luminosity--size--velocity dispersion--metallicity distribution for more distant dwarf spheroidal galaxies in the Local Group. Thus we see no indication that Sagittarius is in any way anomalous.  Further work on the debris streams of Sagittarius is needed to constrain better its initial total luminosity, and to distinguish between purely stellar or dark-matter dominated progenitors.  It is certainly encouraging that our models could reproduce the data available both on the main body and on the debris streams. We wish to stress however, that this does not mean that we have found the ``ultimate'' model. Other models with similar characteristics may also exist. Alternatives would include progenitors with smaller stellar masses or larger dark halos; flattened systems or with anisotropic velocity distributions; or systems with a stellar disk and a spherical dark halo (as proposed for the progenitors of dSph by Mayer et al. (2000)). Moreover, our assumption of a rigid Galactic potential, which does not vary in time over 12 Gyr, is clearly simplistic in view of current models for the formation of structure in the Universe. Only when we have a better estimate of the total luminosity of Sagittarius, both in its main body, as well as on its streams, we will be able to model it in greater detail. The present interest in the debris streams of Sagittarius will help us understand not only the properties of what has turned out to be just another dwarf spheroidal, but also the formation history of our Galaxy. A complete map of the streams will, for example, allow us to derive the Galactic potential (Johnston et al. 1999). If these streams are less smooth or broader than expected, this may indicate smaller scale structure present in the halo either now or when this was assembled.  \\vspace*{-0.25cm}"
    },
    "0002/astro-ph0002211_arXiv.txt": {
        "abstract": "{We quantify the energy radiated through all the collisionally excited lines in a photoionized nebula which is permeated by temperature fluctuations. We assume that these correspond to hot spots which are the results of an unknown heating process distinct from the photoelectric heating.  We consider all the effects of using a higher mean temperature (as compared to the equilibrium temperature) due to the fluctuations not only on each emission line but also on the ionization state of the gas.  If this yet unknown process was to radiate a fixed amount of energy, we find that the fluctuations should correlate with metallicity $Z$ when it exceeds 0.7 solar. The excess energy radiated in the lines as a result of the fluctuations is found to scale proportionally to their amplitude \\tsq. When referred to the total energy absorbed through photoionization, the excess energy is comparable in magnitude to \\tsq.}   \\nonstopmode  \\begin{document} ",
        "introduction": "\\label{sec:intro}  The temperatures of photoionized nebulae are observed to be significantly lower when derived using  recombination lines rather than from forbidden line ratios (e.g. Peimbert et~al. 1995).  This phenomenon has been ascribed to the existence of temperature fluctuations permeating the nebulae. Assuming that the  fluctuations inferred by various authors (e.g. Peimbert et~al. 1995, Esteban et~al. 1998, Rola \\& Stasi\\'nska 1994) are caused by an additional albeit {\\it unknown} heating agent (beside photoionization), we proceed to quantify the energy contribution which this unknown heating process must contribute to the total energy budget of the nebula in order to account for the much higher temperatures characterizing the collisionally excited lines as compared to those inferred from recombination lines (or nebular Balmer continuum). In order to study the effects of arbitrary temperature fluctuations, we first describe the modifications made to the multipurpose photoionization-shock code \\map\\ (Ferruit et~al. 1997). In   Section\\,\\ref{sec:cal}, we present photoionization calculations in which we consider different levels of fluctuation amplitudes and quantify how they alter the global energy budget of the nebula. A brief discussion is presented in Section\\,\\ref{sec:con}. ",
        "conclusions": "\\label{sec:con}  If the turbulences are the result of heating by a yet undiscovered process, we infer from Fig.~\\ref{fig:tsqh} and \\ref{fig:tsqa} that the extra energy radiated via {\\it all} the collisionally excited lines due to the turbulences is a substantial fraction ($\\gama >10$\\%) of the total photoionization energy whenever $\\tsq\\ \\ga 0.02$ for planetary nebulae and $\\tsq\\ \\ga 0.08$ for \\hii\\ regions, respectively.  Taking instead as reference only the heating by photoionization (\\gama), these limits are reduced by two. In this case almost all nebulae would radiate more than 10\\% (\\gama) of their energy as a result of $T$ fluctuations. In effect, typical empirically determined \\tsq\\ values for galactic \\hii\\ regions lie in the range 0.02--0.06 (Luridiana 1999). For extragalactic \\hii\\ regions (Luridiana et~al. 1999, Gonz\\'alez-Delgado et~al. 1994) and planetary nebulae (Peimbert et~al. 1995), even larger values have been encountered, of order 0.1 or more\\footnote{Measured values of \\tsq\\ in excess of 0.1 should be considered only as rough estimates since the assumed regime of small fluctuations does not apply anymore.}. In those cases, the energy envolved can be a large fraction of the energy budget especially if the exciting stellar temperature exceeds $10^5\\,$K.  The underlying assumption behind our calculations is that an external heating agent is at work to generate $T$ fluctuations. We should point out that there exist alternative explanations to the fluctuations which rest on photoionization alone and have not been completely ruled out. Possible mechanisms which we plan to study in some details:  \\begin{enumerate} \\item{\\it Metallicity inhomogeneities} Temperature fluctuations would result naturally from nebulae consisting of a multitude of condensations of greatly varying metallicities (see Torres-Peimbert, Peimbert \\& Pe\\~na 1990) \\item{\\it Gas expansion} Fluctuations could arise from an outflowing wind generated at the surface of dense condensations (problyds) as a result of the champagne effect.  Rapid adiabatic expansion would result in overcooled emission regions in the wake of the wind (resulting in negative $T$ fluctuations with respect to \\teq). \\item{\\it Ionizing source variability} A rapidly varying ionizing field will generate a sequence of outwardly propagating ionization and recombination fronts. The basic asymmetry existing between ionization fronts (propagating at a speed limited by changes in opacity) and recombination fronts (non-propagating), would lead to temperature fluctuations. \\end{enumerate}  To our knowledge, none of these mechanisms have yet been incorporated explicitly in any nebular model and we therefore cannot assess their particular merits. Mechanisms which are good candidates to explain the extra heating assumed in this work (see also discussion by Stasi\\'nska 1998) and which are not directly related to photoionization, have been proposed and include the following processes: \\begin{enumerate} \\item{\\it Shock heating} Either as a result of stellar wind or from supernovae. In the case of giant \\hii\\ regions, this later mechanism was   favored by Luridiana et~al. (1999). \\item{\\it Magnetic heating} Turbulent dissipation from Alf\\'en waves and magnetic line reconnection has been proposed to explain the variations of the \\nii/\\ha\\ ratio emitted by the warm diffuse ionized gas in the Galaxy (Reynolds et~al. 2000). In the case of planetaries, the large values of \\tsq\\ found would then make sense in the light of the success of the magnetically accelerated wind models of Garc\\'\\i a-Segura (2000).  Furthermore, Peimbert et~al. (1995) has shown that the nebulae with the highest gas velocity dispersion show  the highest values of \\tsq, which would be consistent with an increasing role played by magnetic acceleration of the gas in the objects with the largest velocity dispersion. \\end{enumerate}  In summary, the current work shows how \\gama\\ and \\gamh\\ vary with metallicity and \\tsq\\ as well as how these quantities are affected by the ionizing energy distribution. Our results can be translated into energy requirements (function of \\tsq) which  competing explanations of the temperature fluctuations must satisfy and, hence, can be used to probe their respective viability."
    },
    "0002/astro-ph0002094_arXiv.txt": {
        "abstract": "We report the detection of water ice in the Centaur 2060 Chiron, based on near-infrared spectra (1.0 - 2.5 $\\micron$) taken with the 3.8-meter United Kingdom Infrared Telescope (UKIRT) and the 10-meter Keck Telescope.  The appearance of this ice is correlated with the recent decline in Chiron's cometary activity: the decrease in the coma cross-section allows previously hidden solid-state surface features to be seen.  We predict that water ice is ubiquitous among Centaurs and Kuiper Belt objects, but its surface coverage varies from object to object, and thus determines its detectability and the occurrence of cometary activity. ",
        "introduction": "The Centaurs are a set of solar system objects whose orbits are confined between those of Jupiter and Neptune.  Their planet-crossing orbits imply a short dynamical lifetime ($10^6 - 10^7$ yr).  The current belief is that Centaurs are objects scattered from the Kuiper Belt that may eventually end up in the inner solar system as short-period comets.  The first discovered and brightest known Centaur, 2060 Chiron, is relatively well studied.  The object is firmly established as a comet, with a weak but persistent coma.  It is well documented that Chiron possesses neutral colors (e.g., Hartmann et al. 1990, Luu and Jewitt 1990), and a low albedo of $0.14^{+0.06}_{-0.03}$ (Campins et al. 1994).  (It must be noted that most of these measurements were made when Chiron clearly exhibited a coma so that the measurements are likely to have been contaminated by dust scattering from the coma.  The albedo, in particular, should be viewed as an upper limit).  Chiron has a rotation period of $\\sim 6$ hr (e.g., Bus et al. 1989) and a photometric amplitude that is modulated by the cometary activity level (Luu and Jewitt 1990, Marcialis and Buratti 1993, Lazzaro et al. 1997).  Published optical and near-IR spectra of Chiron show a nearly solar spectrum, varying from slightly blue to completely neutral (Hartmann et al. 1990, Luu 1993, Luu et al. 1994, Davies et al. 1998), and devoid of specific mineralogical features.  As a group, the Centaurs display remarkable spectral diversity (e.g., Luu and Jewitt 1996).  The Centaur 5145 Pholus is among the reddest bodies in the solar system (Mueller et al. 1992, Fink et al. 1992), and shows absorption features at 2.00 and 2.25 $\\micron$ (see, e.g., Davies et al. 1993a, Luu et al. 1994).  Cruikshank et al. (1998) interpreted the 2.0 $\\micron$ feature as due to water ice, and the 2.27 $\\micron$ due to methanol.  They derived a best fit to the spectrum which consisted of carbon black and an olivine-tholin-water-methanol mixture. Pholus's extreme red color and low albedo ($0.044 \\pm 0.013$, Davies et al. 1993b) strongly suggest that long-term irradiation of carbon- and nitrogen-bearing ices has resulted in an organic-rich, dark \"irradiation mantle\" (e.g., Johnson et al. 1987).  The spectral differences between Chiron and Pholus have been attributed to the presence of cometary activity in Chiron and the lack thereof in Pholus (Luu 1993, Luu et al. 1994).  A continuous rain of sub-orbital cometary debris falling onto the surface of Chiron may have buried a more primordial irradiation mantle with unirradiated matter ejected from the interior.  In this paper we show further evidence supporting this hypothesis, and that scattering effects by coma dust particles also eliminate spectral features seen in solid surface reflection. ",
        "conclusions": "\\subsection{The spectra}  The Chiron spectra (Fig. 1) show that: (a) Chiron is nearly neutral in the 1.0 -- 2.5 $\\micron$ region; (b) there is a subtle but definite absorption feature at 2 $\\micron$ ($\\sim$ 0.35 $\\micron$ wide, $\\sim 10\\%$ deep) in spectra from 1996 and 1999, and a marginal absorption feature near 1.5 $\\micron$ in the 1999 spectrum; and (c) the spectral slope and the strength of the 2 $\\micron$ feature change with time.  Reflectivity gradients in the JHK region span the range $S'$ = -2 \\%/1000~\\AA\\ to $S' =$ 1\\%/1000~\\AA\\ (Table 2).  In Fig. 2 we compare Chiron spectra from 1993 (from Luu et al. 1994) with the present observations.  The flat and featureless 2 $\\micron$ spectrum from 1993 stands in sharp contrast with the later spectra.  The presence of the 2 $\\micron$ feature in different spectra taken with different telescopes, instruments, and spectral resolutions provides convincing evidence that it is real.  The 2 $\\micron$ and 1.5 $\\micron$ features are clear signatures of water ice, and the shallowness of the features (compared to that of pure water ice) indicates that this ice is mixed with dark impurities (see Clark and Lucey 1984).  We note that the Chiron spectra are very similar to spectra of minerals and water ice (compare the spectra in Fig. 2 with Fig. 14 and 15 of Clark 1981, respectively). In Fig. 3 we show that the Keck Chiron spectrum is well fitted by a model consisting of a linear superposition of a water ice spectrum and an olivine spectrum.  The olivine spectrum is needed to provide the required continuum slope, and at the very small grain size used, the spectrum of olivine is essentially featureless.  The water ice spectrum was calculated based on the Hapke theory for diffuse reflectance (Hapke 1993) and used a grain diameter of 1 $\\micron$.  A description can be found in Roush (1994).  However, we caution that the model is non-unique, and due to the many free parameters in the model (e.g., grain albedo, porosity, roughness), the $1 \\micron$ grain size should not be taken literally.  Similarly, olivine could also be replaced in the fit by other moderately red, featureless absorbers.  The 2 $\\micron$ feature in Chiron is clearly time-variable: it was not apparent in 1993 but changed to a depth of 8 - 10\\% in 1996 and 1999.  Chiron's lightcurve variations can be explained by the dilution of the lightcurve by an optically thin coma (Luu and Jewitt 1990).  Based on this model, we estimate that the 1993 coma cross-section was $\\sim 1.5$ times larger than in 1996 or now. The likeliest explanation for the time-variability of the $2 \\micron$ feature is the degree of cometary activity in Chiron.  In 1993, Chiron's activity level was high (Luu and Jewitt 1993, Lazzaro et al. 1997), resulting in a featureless spectrum dominated by scattering from the coma. By 1996, when the UKIRT observations were made, Chiron's total brightness had dropped by $\\sim 1$ mag to a minimum level (comparable to that of 1983--1985, Lazzaro et al. 1997), leaving spectral contamination by dust at a minimum.  \\subsection{Implications of water ice on Chiron}  \\subsubsection{Cometary activity in Centaurs}  Considering (1) Chiron's time-variable spectrum, (2) the presence of surface water ice, and (3) Chiron's persistent cometary activity, we conclude that Chiron's surface coverage is {\\it not} dominated by an irradiated mantle but more probably by a layer of cometary debris.  Occultation observations suggest that sublimation by supervolatiles (e.g., CO, N$_2$) on Chiron occurs in a few localized icy areas (Bus et al. 1996). Dust grains ejected at speeds $< 100$ m s$^{-1}$ (the escape velocity) will reimpact the surface, building a refractory layer which tends to quench sublimation. Nevertheless, the outgassing is still sustained by the sporadic exposure of fresh ice on the surface.  Sublimation experiments with cometary analogs illustrate this phenomenon: outgassing produces a dust layer, but fresh icy material can still be periodically exposed by avalanches and new vents created by impacts from large dust particles (Gr\\\"un et al.  1993). If, in keeping with a Kuiper Belt origin, Chiron once possessed an irradiation mantle, we suspect that it has been either blown off by sublimation or buried under a dust layer thick enough to mask its features.  If so, the present low albedo of Chiron would be due to cometary dust particles rather than irradiated material.  In contrast, the lack of cometary activity in the Centaur Pholus is consistent with an encompassing surface coverage by the irradiation mantle -- witness the extreme red color and absorption features associated with hydrocarbon materials.  Although water ice may exist locally on the surface (Cruikshank et al. 1998), Pholus's spectral properties are still dominated by the organic irradiated crust. If this hypothesis is correct, cometary activity should be uncommon among Centaurs with very red colors (irradiated material), and more common among those with neutral (ice) colors.  As the observational sample of Centaurs grows, this is a simple prediction that can be directly tested. However, it remains unclear why cometary activity was activated on Chiron and not on Pholus, even though the two Centaurs are at similar heliocentric distances.  One possibility is that Pholus was recently expelled from the Kuiper Belt and has not yet been heated internally to a degree sufficient to blow off the irradiation mantle, but this hypothesis cannot be easily tested, given the chaotic nature of the orbits of Centaurs.  \\subsubsection{Centaur and Kuiper Belt surfaces}  We summarize the spectral properties of Centaurs and KBOs in Table 3. Thus far, water ice has been reported in 3 Centaurs (Chiron, Pholus, 1997 CU$_{26}$) and 1 KBO (1996 TO$_{66}$).  The existing data are too sparse to establish whether a correlation exists between color and the abundance of water ice among Centaurs and KBOs. However, water ice is present in all three Centaurs for which near-IR spectra are available, and in 1 out of 3 studied KBOs.  The preponderance of water ice among Centaurs makes us suspect that that the \"low\" rate of detection of water ice in KBOs has more to do with the faintness of the targets and the resulting low-quality spectra than with the intrinsic water contents in KBOs.  Considering the existing data and the high cosmochemical abundance of water ice, we predict that water ice is ubiquitous among {\\it all} objects that originated in the Kuiper Belt, although the amount might vary from one object to another and thus determines the possibility for cometary activity in these bodies.  In short, it would be a good idea to re-observe those KBOs which show no apparent water ice feature (1993 SC and 1996 TL$_{66}$) at higher signal-to-noise ratios.  Water ice might be present after all."
    },
    "0002/astro-ph0002431_arXiv.txt": {
        "abstract": "We fit the SEDs of the TeV blazar Mkn 501 adopting the homogeneous Synchrotron-Self Compton model to simultaneous X-ray and TeV spectra recently become available. We present detailed model spectra calculated with the above constraints and taking into account the absorption of TeV photons by the IR background. We found that the curved TeV spectra can be naturally reproduced even without IRB absorption. Taking IRB absorption into account changes the required parameter values only slightly. ",
        "introduction": "The study of Blazars has been recently enriched by the detection of very high energy gamma rays (energy in the TeV range) from a handful of nearby sources. This discovery opens the possibility to effectively test and constrain the radiative mechanisms invoked to explain the emission from Blazars and to investigate the physical conditions in relativistic jets.  As discussed in Tavecchio et al. 1998 (hereafter T98), the knowledge of the simultaneous X-ray and TeV spectra of Blazars allows to univocally find the set of physical parameters necessary within  a homogeneous Synchrotron--Self Compton model with an electron distribution described by a broken power-law. Here we apply the SSC model to recent high quality, simultaneous X-ray and TeV data of the well studied source Mkn 501 and discuss the results. ",
        "conclusions": "The parameters adopted for the model are reported in Table 1. Our model indicates a rather low value for the magnetic field and a high $\\gamma _{break}$ (see Tab.1). The transition from the low state to the high state in Mkn 501 is consistent with an increase of a factor of 2 in both $\\gamma _{break}$ and B and with an almost constant value of $\\delta $ and $R$. These results are in agreement with a similar analysis of the April 16 flare by Bednarek \\& Protheroe (1999): for a minimum timescale of 2.5 h they found that the TeV spectrum can be well reproduced by $B\\sim 0.03$ and $\\delta \\sim 15$.  Although some authors (see Konopelko et al. 1999) have recently proposed that the curvature of the TeV spectrum of Mkn 501 provides evidence for absorption by IRB, the curvature could be intrinsic and related to the curved electron distribution necessary to fit the X-ray data. Our models show that a curved spectrum is a plausible explanation for the curvature. The introduction of the IRB does not dramatically change the inferred physical parameters. As shown in Fig.2 IRB/no IRB TeV spectra are very similar up to 15 TeV, while for higher energy the power-law spectrum of the no IRB case is changed in an exponential profile. Therefore spectral data above 10 TeV are needed in order to understand if IRB absorption affects Blazar spectra.  Finally we note that in the theory of particle acceleration by shocks (for a recent critical discussion see Henri et al. 1999) the maximum Lorentz factor of accelerated electrons, obtained equating acceleration time and cooling time, is given by $\\gamma _{max}\\simeq 10^7 (B/0.02 G)^{-1/2}$, where we used the value of B we found in the high state of Mkn 501. Therefore our fit suggests that during the high activity states of Mkn 501 the electrons can reach the maximum energy fixed by the balance between cooling and acceleration processes. In states of lower activity the acceleration time could be longer and the condition $t_{esc}<t_{acc}$ (where $t_{esc}$ and $t_{acc}$ are the escape time and the acceleration time respectively) could prevent the electrons for reaching the maximum energy.  \\begin{figure} \\vskip -4 cm \\hskip 0.5 truecm \\centerline{\\psfig{figure=due_sed.ps,width=11 truecm,clip=,angle=-90}} \\vspace*{-3.0cm} \\caption{TeV spectrum of Mkn 501-High state without (left) and with (right) IRB absorption. We also show the contribution from electrons with different Lorentz factors. From the left the curves show the emission from electrons with Lorentz factor in the range: $1-10^6$, $10^6-3\\times10^6$, $3\\times10^6-10^7$, $10^7-3\\times10^7$, $3\\times10^7-10^8$} \\end{figure}"
    },
    "0002/astro-ph0002425_arXiv.txt": {
        "abstract": "The formation rate of luminous galaxies seems to be roughly constant from $z\\sim 2$ to $\\sim 4$ from the recent observations of Lyman break galaxies (LBGs) (Steidel \\etal 1999). The abundance of luminous quasars, on the other hand, appears to drop off by a factor of more than twenty from $z\\sim 2$ to $z\\sim 5$ (Warren, Hewett, \\& Osmer 1994; Schmidt, Schneider, \\& Gunn 1995). The difference in evolution between these two classes of objects in the overlapping, observed redshift range, $z=2-4$, can be explained naturally, if we assume that quasar activity is triggered by mergers of luminous LBGs and one quasar lifetime is  $\\sim 10^{7-8}~$yrs. If this merger scenario holds at higher redshift, for the evolutions of these two classes of objects to be consistent at $z>4$, the formation rate of luminous LBGs is expected to drop off at least as rapidly as $\\exp\\left(-(z-4)^{6/5}\\right)$ at $z>4$. ",
        "introduction": "Observations of galaxies in the rest frame UV band (\\cite{Lilly96}1996; \\cite{Madau96}1996; \\cite{Connolly97}1997; \\cite{Sawicki97}~1997; \\cite{Treyer98}1998; \\cite{Pascarelle98}~1998) indicate that the galaxy formation rate rises steeply from $z=0$ to $z\\sim 1$, with a nearly constant rate thereafter up to $z\\sim 4$ % (\\cite{Steidel99}1999). While at low redshift ($z<2$) the evolution of luminous quasar abundance resembles that of luminous galaxies (e.g., Sanders \\& Mirabel 1996; Boyle \\& Terlevich 1998), at high redshift ($z>2$) the two classes of objects do not seem to parallel one another, with the luminous quasar formation rate (e.g., Warren \\etal 1994; Schmidt \\etal 1995) dropping off more steeply than that of luminous galaxies.   In this {\\it Letter} a phenomenological approach is taken to relate the {\\it observed} formation rate of luminous LBGs to the {\\it observed} abundance evolution of luminous quasars at $z>2$. It is shown that, {\\it if 1) quasar activity is triggered by LBG mergers} and {\\it 2) each quasar period lasts $\\sim 10^{7-8}$yrs}, then the apparent difference, {\\it in both shape and amplitude}, between the evolutions of bright LBGs and bright quasars from $z=2$ to $z=4$ can be explained quantitatively. The first assumption finds its support from both the observational evidence that a significant fraction of quasar hosts have disturbed morphologies or ongoing galaxy-galaxy interactions (e.g., Boyce \\etal 1996; Bahcall \\etal 1997; Boyce, Disney, \\& Bleaken 1999) and the theoretical consideration that merger of two (spiral) galaxies seems to provide a natural mechanism to fuel the central black hole (e.g., Barnes \\& Hernquist 1991). The second assumption is also theoretically well motivated (Rees 1984, 1990) and now strongly implied or required by the mounting observational evidence that most nearby massive galaxies seem to harbor inactive black holes at their centers (e.g., Richstone \\etal 1998).  The primary purpose of this work is to use this merger model to infer the LBG formation rate at higher redshift $z>4$. Given the precipitous drop-off of luminous quasar ($M_B<-26.0$) abundance from $z=2$ to $z=5$ the formation rate of luminous ($M_{AB}\\geq -23$ to $-22$) LBGs at higher redshift ($z>4$) is predicted to drop off as least as fast as $\\propto \\exp\\left(-(z-4)^{6/5}\\right)$, if the merger scenario holds. A cosmological model with $q_0=0.5$ and Hubble constant $H_0=50$km/sec/Mpc is assumed for the analysis presented here.  It is noted that this simple merger model would probably fail at $z<2$ without having taken into account the evolution of gaseous fuel supply to the central black holes in galaxies (Kauffmann \\& Haehnelt 1999; Haiman \\& Menou 1999). ",
        "conclusions": "In an early classic paper Efstathiou \\& Rees (1988) show that quasar abundance at high redshift can be accounted for in the standard cold dark matter model {\\it if massive halos are related to the formation of black holes}, with an intriguing prediction that the abundance of luminous quasars should decrease rapidly beyond $z=5$ (for a more recent treatment see Haehnelt \\& Rees 1993). [The evolution of low-luminosity quasars, of course, does not necessarily have to follow that of their luminous counterparts (e.g., Haiman \\& Loeb 1998)].  In this {\\it Letter} a different approach is taken by directly relating the {\\it observed} evolution of luminous LBGs to the {\\it observed} evolution of luminous quasars at high redshift ($z>2$). With a set of seemingly reasonable parameter values, it is shown that consistency between the two classes of objects at $z=2-4$, where both classes are observed, can be achieved, if one assumes that 1) {\\it Lyman break galaxies merger to trigger quasar activity} and 2) {\\it quasar lifetime is $\\sim 10^{7-8}$yrs}. At $z>4$, consistency can be achieved, {\\it only if additionally} the formation rate of luminous LBGs drops off as $\\exp(-(z-4)^{6/5})$ or faster, a prediction that may be tested by future observations. One implication from this model is that LBGs with $M_{AB}\\geq -23$ to $-22$ merge to form quasars with  $M_B<-26.0$ at $z>2$. Correlation analysis of relevant LBGs and quasars should shed light on this.   At lower redshift additional, more model dependent assumptions regarding the supply of available gas to fuel black holes would be required to make qualitatively viable predictions. Kauffmann \\& Haehnelt (1999; see also Haiman \\& Menou 1999) have presented a detailed model, based also on merger scenario, to unify the evolution of galaxies and quasars in the cold dark matter model under several plausible assumptions concerning the evolution of fuel gas to the central black holes. The success of the model of Kauffmann \\& Haehnelt (1999) at low redshift ($z<2$) and the model presented here at high redshift ($z>2$), both based on galaxy merger scenario, suggests that {\\it galaxy merger may play an indispensable role in quasar formation.}"
    },
    "0002/astro-ph0002339_arXiv.txt": {
        "abstract": "Line-driven winds from hot stars and accretion disks are thought to adopt a unique, critical solution which corresponds to maximum mass loss rate and a particular velocity law. We show that in the presence of negative velocity gradients, radiative-acoustic (Abbott) waves can drive shallow wind solutions towards larger velocities and mass loss rates. Perturbations introduced downstream from the wind critical point lead to convergence towards the critical solution. By contrast, low-lying perturbations cause evolution towards a mass-overloaded solution, developing a broad deceleration region in the wind. Such a wind differs fundamentally from the critical solution. For sufficiently deep-seated perturbations, overloaded solutions become time-dependent and develop shocks and shells. ",
        "introduction": "Atmospheres of hot luminous stars and accretion disks in active galactic nuclei and cataclysmic variables form extensive outflows due to super-Eddington radiation fluxes in UV resonance and subordinate lines. An understanding of these winds is hampered by the pathological dependence of the driving force on the flow velocity gradient. Castor, Abbott, \\& Klein (1975; CAK hereafter) found that line-driven winds (hereafter LDWs) from O~stars should adopt a unique, critical state which corresponds to maximum mass loss rate. The equation of motion for a 1-D, spherically symmetric, polytropic outflow subject to a Sobolev line force allows for two infinite families of so-called shallow and steep solutions. However, none of these families can provide for a {\\it global} solution alone. Shallow solutions do not reach infinity, while steep solutions do not extend into the subsonic regime including the photosphere. The critical wind starts then as the fastest shallow solution and switches at the critical point in a continuous and differentiable manner to the slowest steep solution. Hence the critical point and not the sonic point determines the bottleneck in the wind. This description in principle applies equally to winds from stars and accretion disks.  A physical interpretation of the CAK critical point was given by Abbott (1980), who derived a new type of radiative-acoustic waves (hereafter Abbott waves). These waves can propagate inward, in the stellar rest frame, only from below the CAK critical point. Above the critical point, they are advected outwards. Hence, the CAK critical point serves as an information barrier, much as the sonic or Alfv\\'en points in thermal and hydromagnetic winds. Abbott's analysis was challenged by Owocki \\& Rybicki (1986) who found for a pure absorption LDW the signal speed to be the sound speed and not the much faster Abbott speed. As noted already by these authors, this should be a consequence of assuming pure line absorption, which does not allow for any radiatively modified, inward wave mode. Meanwhile there is ample evidence for Abbott waves in time-dependent wind simulations (Owocki \\& Puls 1999).  Shallow solutions fail to reach infinity because they cannot perform the required spherical expansion work, implying that the flow starts to decelerate. Since this usually occurs very far out in the wind, the local wind speed is much larger than the local escape speed, and the wind escapes to infinity. Thus, a simple generalization of the CAK model allowing for flow deceleration renders shallow solutions globally admissible. This raises a fundamental question of why the wind would adopt the critical solution at all, and attain the critical mass loss rate and velocity law, as proposed by CAK.  In this Letter we analyze a physical mechanism which drives shallow solutions towards the critical one, and discuss under what conditions this evolutions does not terminate at the CAK solution, but continues into the realm of overloaded solutions. We find that simulations so far were affected by numerical runaway towards the critical solution, by not accounting for Abbott waves in the Courant time step. ",
        "conclusions": "We find that shallow solutions to line-driven winds are subcritical with respect to Abbott waves (sub-abbottic). These waves cause shallow solutions to evolve towards larger speeds and mass loss rates because of the asymmetry of the line force with regard to positive and negative velocity gradients and because perturbations with opposite signs of $dv/dz$ propagate in opposite directions. Steep velocity slopes propagate towards the wind base, steepen the inner wind and lift it to higher mass loss rates.   In the presence of enduring wind perturbations,  this proceeds until a critical point forms and Abbott waves can no longer penetrate inwards.  The resulting solution does not necessarily correspond to the CAK wind. For perturbations which originate below the critical point, the developing Abbott wave barrier is found to be the termination point of a mass-overloaded solution. The velocity law acquires a kink at the termination point, where the wind starts to decelerate. Whether the wind converges to a critical or overloaded solution depends entirely on the {\\it location} of perturbations, and not, e.g., on boundary conditions at the wind base.  If Abbott waves are not accounted for in the Courant time step of hydrodynamic simulations, we find that {\\it numerical} runaway can drive the solution towards the critical CAK wind. A detailed discussion of this will be given elsewhere.  Future work has to clarify whether and where perturbations causing local flow deceleration, $dv/dz<0$, can occur in LDWs. Overloaded winds may be detected observationally. While their mass loss rates should still be close to CAK values, broad regions of decelerating flow could be identified in P~Cygni line profiles. Furthermore, shocks occurring in overloaded solutions with infalling gas may contribute to the X-ray emission from LDWs, besides shocks from the line-driven instability (Lucy 1982; Owocki et al.~1988). Note that the present wind runaway occurs already in the lowest order Sobolev approximation, and is therefore, unrelated to the line-driven instability which depends on velocity curvature terms (Feldmeier 1998)."
    },
    "0002/astro-ph0002080_arXiv.txt": {
        "abstract": "We examine the Fisher--Tully $cz<1000$ km s$^{-1}$ galaxy sample to determine whether it is a complete and representative sample of {\\it all} galaxy types, including low surface brightness populations, as has been recently claimed. We find that the sample is progressively more incomplete for galaxies with (1) smaller physical diameters at a fixed isophote and (2) lower HI masses. This is likely to lead to a significant undercounting of nearby gas-rich low surface brightness galaxies. However, through comparisons to other samples we can understand how the nearby galaxy counts need to be corrected, and we see some indications of environmental effects that probably result from the local high density of galaxies. ",
        "introduction": "Two of the key questions in 20th century extragalactic studies have concerned the density and the composition of the galaxy population in the Universe. Our knowledge of galaxy types and their abundance depends critically on the issue of completeness of our galaxy catalogs. Since our current catalogs are constructed by various observational means, they are, by definition, limited by natural and technological selection effects.  Recently, Briggs (1997a) has argued that the Fisher \\& Tully (1981, F--T) catalog of nearby galaxies is complete to its redshift and sensitivity limits, even for low surface brightness (LSB) galaxies. F--T examined HI emission from a sample of 1787 angularly large galaxies accessible from the Green Bank radio telescopes. They believed the sample to be very complete for late-type galaxies within a redshift of $cz=1000$ km s$^{-1}$, with angular diameters larger than 3$'$, located in regions at $|b|>30^\\circ$ and $\\delta>-33^\\circ$. Briggs additionally found a sensitivity limit depending on the HI mass and defined the F--T ``completeness zone'' as extending out to: \\begin{equation} z_{CZ} = \\left\\{ \\begin{array}{ll} 1000\\mbox{km s}^{-1}/c & \\mbox{if $M_{HI}>10^{8.45} h_{75}^{-2} M_\\odot$} \\\\ (M_{HI}h_{75}^{2}/10^{8.45}M_\\odot)^{5/12} 1000\\mbox{km s}^{-1}/c & \\mbox{otherwise.} \\end{array} \\right. \\end{equation} The form of the limit for sources with masses smaller than $10^{8.45}M_\\odot$\\footnote{We adopt $H_0=75$ km s$^{-1}$ Mpc$^{-1}$ for masses and distances quoted hereafter, and we use the heliocentric velocity corrected by 300 km s$^{-1} \\cos b \\sin l$ as in F--T and Briggs.} was based on a semi-analytic, semi-empirical fit to the HI sensitivity.  Briggs pointed out that surveys of low surface brightness galaxies (Schneider et al.~1990; Schombert et al.~1992; Matthews \\& Gallagher 1996; Impey et al.~1996) have identified sources that are primarily at larger distances, but they have added very few within the F--T ``completeness zone.'' He concluded that LSB galaxies ``must already be fairly represented by nearby, previously cataloged examples.''  This is an interesting idea, but the conclusion does not necessarily follow from the analysis for several reasons: (1) For sources with low HI masses $z_{CZ}$ is so small that redshift distances are very uncertain and Galactic HI emission creates strong confusion. A further objection to Briggs' analysis is that (2) the F--T sources are themselves incomplete within Briggs' ``completeness zone'' because of angular size selection effects.  Lastly, (3) the $cz<1000$ km s$^{-1}$ region around the Local Group has about twice the average galactic density and therefore is not a very representive sample of the Universe as a whole.  In a companion paper, Briggs (1997b) used the F--T and LSB samples to derive an HI mass function. He commented on the need for corrections for incompleteness and noted the usefulness of a  ${\\cal V}/{\\cal V}_{max}$ test for establishing whether the LSB samples were complete, but he did not discuss the problems the F--T sample has in this regard. In this paper we explore the limitations of the F--T sample and discuss how it and more recent surveys may be properly used to understand the composition of galaxy populations. We show that the local samples of galaxies display morphological segregation characteristics associated with high density environments. Finally, we find that in all HI mass ranges the dominant class of galaxies are those with the smallest angular diameters at the isophotal limits of the original Palomar Sky Survey (PSS-I). These galaxies may be physically small or appear small because they are LSB systems; in either case, they are greatly under-represented in the F--T sample and most other optical surveys. ",
        "conclusions": "Optically selected samples favor optically bright galaxies. This truism holds even for diameter-limited galaxy surveys because LSB galaxies appear small at the surface-brightness limit of the optical images (Disney 1976). The HI-Slice survey (Spitzak \\& Schneider 1998), which is unbiased by galaxy diameter, indicates that the optically smallest galaxies are the most common. Current searches on deep photographic plates for small angular diameter sources (Eder \\& Schombert 1999) are also uncovering indications of this population. Such LSB and 21 cm surveys are successfully probing sources with HI masses down to $\\sim10^8M_\\odot$, but for lower mass objects HI flux and angular size limitations of existing surveys allow detections of these sources to only a few Mpc.  Because of the small distances at which low mass and LSB sources are accessible, we need to consider the possible impact of the local environment on them. Although the large scale distribution of LSB galaxies appears similar to that of other galaxies (Mo, McGaugh and Bothun 1994), on scales of less than 2 Mpc their numbers drop off sharply. The most likely explanations are that either LSB disks are fragile and easily disturbed by other galaxies, or tidal interactions induce star formation that converts LSB galaxies into normal Hubble type objects. Regardless of the underlying cause why LSB galaxies avoid high density regions, this fact produces an expectation that the local region of space will be deficient in the number of LSB disk galaxies due to the large number of high mass spirals and the proximity to the very dense Virgo Cluster.  Other influences of the local environment may also play a role in the distribution of galaxy types. In high density regions, galaxies are often gas deficient for their morphological type because of stripping or evaporation, and morphological segregation favors earlier-type, less gas rich galaxies. Both of these effects would tend to {\\it lower} the percentage of gas-rich systems nearby.  Briggs (1997a) asked ``Where are the nearby gas-rich low surface brightness galaxies?'' The answer depends on the mass range of objects being studied. There appears to be a local deficit of high-mass LSB systems, which is probably an environmental effect. The story for low-mass systems is less settled, because of the difficulty in detecting them to any significant distance, but it is clear there is a much larger population of small-optical-diameter galaxies than optical surveys have previously revealed. And finally, since these low-mass objects have not yet been detected beyond the local high-density environment, it is possible that they are even more profuse than they appear locally."
    },
    "0002/astro-ph0002049_arXiv.txt": {
        "abstract": "We calculate the spectrum resulting from the interaction of a fireball with ambient soft photons. These photons are assumed to be produced by the walls of a funnel in a massive star. By parameterizing the radial dependence of the funnel temperature we calculate the deceleration of the fireball self--consistently, taking into account the absorption of high energy $\\gamma$--rays due to interaction with the softer ambient photons. The resulting spectrum is peaked at energies in agreement with observations, has a $\\nu^2$ slope in the X--ray band and a steep power--law high energy tail. ",
        "introduction": "We have recently proposed (Lazzati et al. 2000, hereafter Paper I) that the gamma--ray burst (GRB) phenomenon originates from the interaction of a relativistic fireball with a dense photon environment, leading to Compton drag. On one hand this is an inevitable effect if the progenitors of GRBs are massive stars which are about to explode or have just exploded as supernovae; on the other hand this mechanism greatly alleviates the efficiency problem faced by the standard internal shock scenario (Lazzati, Ghisellini \\& Celotti 1999; Panaitescu, Spada \\& Meszaros 1999; Kumar 1999). In Paper I we have discussed the basic Compton drag scenario, showing how this process can convert bulk motion energy directly into radiation with a remarkable high efficiency and, on the basis of simple estimates, how the resulting spectrum should peak, in a $\\nu F(\\nu)$ representation, around $\\sim$1 MeV, as observed.  Here we quantitatively and self--consistently estimate the predicted spectrum, assuming that the fireball propagates in a funnel inside a massive star, and show that, independently of the details of the model, it satisfactorily resembles what observed. Since the funnel walls emit a blackbody spectrum and the scattered photons are boosted by the square of the Lorentz factor ($\\Gamma$) of the fireball, the local spectrum has a blackbody shape, at a temperature enhanced by $\\Gamma^2$. However, the observed spectrum, convolution of all the locally emitted spectra, is not a blackbody, due to four main effects: i) the funnel walls would not be at a uniform temperature, but there should be a gradient between the internal and external parts; ii) if the Compton drag process is efficient, the fireball decelerates; iii) the very high energy emission produced in the internal regions can interact with the ambient photons, producing electron--positron pairs; iv) the fireball may become optically thin to scattering outside the funnel, where the ambient photons are characterized by the same temperature, but their energy density is progressively diluted with distance. ",
        "conclusions": "If the fireball propagates in a dense photon environment the Compton drag effect must necessarily be taken into account, and it may even be the dominant emission mechanism, able to decelerate the fireball without the need of internal shocks and without invoking the build--up of large magnetic fields.  In this letter  we have shown that the predicted spectrum, rather than being simply a black body spectrum boosted in energy, has a complex shape, with power law segments corresponding to the decrease in temperature of the funnel, deceleration of the fireball, and dilution of the radiation energy density as the fireball propagates outside the funnel while remaining optically thick.  The general features of the predicted spectrum qualitatively agree with observations, since they can explain the steep power law high energy tail, the peak of the emission, and a hard tail in the X--ray band. The latter feature is particularly interesting, since other models made different predictions. In the standard internal shock synchrotron model, in fact, the spectrum cannot be harder than $\\nu^{1/3}$ in the thin part, and it is very unlikely that self--absorption can take place in the X--ray band (Granot, Piran \\& Sari 2000). This would in fact imply a huge density of relativistic particles, making the inverse Compton effect largely dominate the total radiation output. This radiation would be emitted at higher and yet unobserved frequencies, and would then worsen the already severe efficiency problem.  In the quasi--thermal Comptonization model, on the other hand, the typical predicted spectral shape in the X--ray band is $\\propto \\nu^0$, down to the typical frequencies of the seed soft photons, i.e. the IR--optical band (Ghisellini \\& Celotti 1999; Meszaros \\& Rees 2000).  The existing observations of a significant fraction of burst spectra harder than $\\nu^{1/3}$ (Preece et al., 1999a,b; Crider et al., 1997) are therefore already a challenge to existing models, and may suggest a Compton drag origin of this portion of the spectrum. However the situation is not already a clear--cut because, to receive enough photons to study the spectral shape, integration times are much longer than the dynamical time--scales of the system, with the spectrum rapidly evolving in time. More sensitive instruments, such as the Burst Alert Telescope (BAT, a coded mask detector more sensitive than BATSE) onboard the foreseen Swift satellite will probably overcome this limitation.  We must also stress that the Compton drag scenario is not alternative to the more conventional internal shock one.  Indeed, the front of the fireball will decelerate first, plausibly causing subsequent un-decelerated parts to shock even if the central engine is working in a continuous way.  This would produce additional radiation, either by the synchrotron and inverse Compton processes or by quasi--thermal Comptonization, depending on the details of the particle acceleration mechanism (see Ghisellini \\& Celotti 1999).  We then expect spectral evolution: since the latter radiation mechanisms produce a steeper low energy tail, a hard--to--soft transition (i.e. from $\\nu^2$ to $\\nu^{1/3}$ or $\\nu^0$) would occur.   In this paper, we have considered the illustrative case of a single fireball moving out through an extended stellar envelope, along a funnel which is empty of matter but pervaded by thermal radiation from the funnel walls. The fireball itself (for typical parameters) remains optically thick until it expands beyond the stellar surface. A burst with complex time-structure could be modeled by a series of fireballs or expanding shells. However, in this more general case, the later shells would suffer less drag, since not enough time may have elapsed to replenish the entire funnel cavity with seed photons. Indeed one expects the spikes to be more powerful the longer is the time interval between them, as more seed photons could pervade the cavity. This, besides causing internal shocks with the first shell which has been efficiently decelerated by Compton drag, will also result in a distribution of $\\Gamma$--factors: they will become greater on axis, where few seed photons can efficiently Compton drag the shells, and smaller towards the border of the funnel, where seed photons can be replenished by the funnel walls.  We plan to investigate these possibilities and their consequences on the associated predicted afterglows in future work."
    },
    "0002/astro-ph0002519_arXiv.txt": {
        "abstract": "We present a detailed analysis of the cluster of galaxies \\a4.  The optical data include a redshift catalogue of 466 galaxies, out of which 274 belong to the main cluster and a CCD photometric catalogue in a much smaller region, with 239 and 610 galaxies in the V and R bands respectively. The X-ray analysis is based on an image obtained with the ROSAT PSPC. \\\\ Besides Abell 496 itself, the velocity distribution along the line of sight shows the existence of at least four structures at different redshifts, one of them seeming to be a poor cluster at a velocity of 30083 km/s. The other of these structures have a too large spatial extent to be clusters but may be filaments along the line of sight or other young structures.\\\\ Various independent methods show that Abell 496 appears to be a quite relaxed cluster, except perhaps for the distribution of emission line galaxies.  These appear to be distributed in two samples falling on to the main cluster, one from the back (the ELGs concentrated towards the west) and one from the front (the high velocity ELGs). \\\\ The bright part of the galaxy luminosity function, built from the redshift catalogue, shows a flattening at R$\\sim 16$ (M$_{\\rm R}\\sim -20.5$), and can be accounted for by a gaussian distribution of bright galaxies and a power law or Schechter function for faint galaxies. The deeper galaxy counts derived from CCD imaging show a dip at R$\\sim 19.5$ (M$_{\\rm R} \\sim -17$) which can be modelled assuming a cut-off in the luminosity function such as that observed in Coma.\\\\ We propose a model for the X-ray gas and derive the galaxy, X-ray gas and total dynamical masses, as well as the baryon fraction in the cluster. \\a4\\ appears as a relaxed cluster which can be used as a prototype for further studies.\\\\ ",
        "introduction": "In the framework of hierarchical clustering, the Universe is believed to be made of galaxies distributed in sheets encircling voids or filaments, at the intersection of which clusters of galaxies are located. Such models can be tested through the analysis of clusters, which are likely to keep a ``memory'' of their formation. This is suggested for example by the alignment effects observed in some clusters, such as for example Abell 3558 (Dantas et al. 1997) or Abell 85 (Durret et al. 1998), where the cD, the brightest galaxies, the X-ray emitting gas and possibly even larger scale structures (in the case of Abell 85) all appear aligned along the same direction. Multi-wavelength studies of clusters of galaxies also allow us to draw a global and coherent portrait of these objects, which we can then use to address other questions of interest, such as the influence of mergers and environmental effects at various scales on the properties of both galaxies and X-ray gas. Large scale (i.e. cluster size) mergers are quite often observed from substructures detected in the X-ray gas; smaller scale mergers (i.e. group size) such as the infall of dwarf galaxies onto groups surrounding bright galaxies can be derived from various methods such as those of Serna \\& Gerbal (1996) or Gurzadyan \\& Mazure (1998), which require optical velocity and magnitude catalogues; the existence of subclustering also has an influence on the shape of the galaxy luminosity function, which in some cases appears to show a deficit of faint galaxies often interpreted as due to accretion of dwarf galaxies onto larger galaxies or groups (e.g. in Coma, Lobo et al. 1997, Adami et al. 2000).  It therefore appears important to analyze cluster properties in detail before using them in other studies.  Note in particular that the existence of substructures can lead to overestimate cluster velocity dispersions, and hence M/L ratios and the value of $\\Omega _0$ in clusters.  With the improvement of both observational means (better X-ray detectors, optical multi-object spectroscopy) and modern methods of analysis (some of which are described below), an ever increasing number of clusters showing evidence for merging and environmental effects has been found. A rather general picture has therefore emerged for clusters, with a main relaxed body on to which groups of various sizes can be falling.  Our approach these last years has been to study a small sample of nearby clusters in detail. These have the advantage of being bright, and can therefore be observed in detail within a reasonable amount of telescope time. Besides, they are free of evolution effects.  We present here a detailed multi-wavelength study of \\a4, based on optical (extensive redshift and photometric catalogues) and X-ray (ROSAT PSPC) data.  \\a4 is a richness class 1 (Abell 1958) cD type (Struble \\& Rood 1987) cluster at a redshift of 0.0331. For a Hubble constant H$_0$=50 km s$^{-1}$ Mpc$^{-1}$, the corresponding scale is 55.0 kpc/arcmin and the distance modulus is 36.52. At optical wavelengths, an adaptive kernel map of the central region (in a 60$\\times$60 arcmin$^2$ square) has revealed a somewhat complicated structure, with a strong concentration of galaxies in the north-south direction (Kriessler \\& Beers 1997). Note however that this map does not include redshift information. Thorough investigations of the X-ray properties of \\a4\\ can be found in Mohr et al. (1999) and Markevitch et al. (1999); their results will be compared to ours in section 4.5.  The paper is organized as follows: we present the data in section 2; the structures along the line of sight derived from the velocity catalogue are described in section 3; the optical properties of the \\a4\\ cluster itself are described in section 4; the X-ray cluster properties are described in section 5; a summary and conclusions are given in section 6. ",
        "conclusions": "The optical analysis of the \\a4\\ field has shown the existence of several structures along the line of sight. Among these, one (structure 6) is likely to be a poor, diffuse and low mass cluster, while two others (structures 4 and 9) are probably filaments more or less aligned along the line of sight, the latter presenting a smooth velocity gradient. Notice that the distances between these various structures are very large.  The cluster \\a4\\ itself has quite a regular morphology. It includes 274 galaxies in the [7813,11860 km/s] velocity range and has a velocity dispersion of 715 km/s. Its velocity distribution implies a small amount of substructure. The analysis of the correlations between position, luminosity and velocity dispersion indicates a post-violent relaxation state. We can notice that both the distance to the cluster center and the overall velocity dispersion ranges are reduced when emission line galaxies (hereafter ELGs) are excluded.  This agrees with the general scheme that ELGs are often found in the outskirts of clusters and are not as strongly tied to the cluster gravitationally (e.g. Biviano et al. 1997).  There may be two samples of ELGs falling on to the main cluster, one from the back (the ELGs concentrated towards the west) and one from the front (the high velocity ELGs).  The bright luminosity function derived from our redshift catalogue shows a flattening at R$\\sim 16$ (M$_{\\rm R}\\sim -20.5$), comparable to similar shapes found in other clusters. This suggests at least a bimodal distribution, one for ellipticals and one for fainter galaxies. The fact that the flattening occurs at the same absolute magnitude as for other clusters suggests that the galaxy populations in all these clusters are comparable.  At fainter magnitudes, galaxy counts derived from CCD imaging show a dip at R$\\sim 19.5$ (M$_{\\rm R} \\sim -17$) which can be reproduced if we assume a magnitude cut-off similar to that observed in Coma (Adami et al. 2000). Notice that such a cut-off is observed in the very central regions of both clusters. Although this result is only based on imaging and remains to be confirmed spectroscopically, we may be evidencing a second example of a cut-off of the faint end of the luminosity function in a cluster.  We have modelled the X-ray gas and derived the X-ray gas mass and the dynamical mass, which we compare to the stellar mass. At the limiting radius (1.62 h$_{50}^{-1}$ Mpc) of the image, we find a fraction of X-ray gas to total mass of 0.12$-$0.15 and a stellar to X-ray gas mass ratio of 0.16.  We can note that \\a4\\ follows exactly the two by two correlations between the X-ray luminosity (L$_{\\rm X}=6.8\\ 10^{44}$ erg/s, Wu et al. 1999), the X-ray temperature (T$_{\\rm X}=4$ keV) and the galaxy velocity dispersion ($\\sigma _v = 715$ km/s) described in the literature (see e.g. Wu et al. 1999 and references therein). These values are typical of a richness class 1 cluster.  \\a4\\ therefore appears to be a relatively quiet and simple cluster, with no strong environmental effects, although we may see an enhancement of the X-ray emission and of the number of emission line galaxies towards the north west.  While Coma has long been the archetype of a relaxed cluster and is not believed to be relaxed any more (see Biviano 1998 and the proceedings of the Coma meeting), the results presented above suggest that \\a4\\ may be such a prototype, and can be used as a ``template'' in the future study of more complex (i.e. substructured) clusters."
    },
    "0002/astro-ph0002033_arXiv.txt": {
        "abstract": "We present multi-frequency VLA, multicolor CCD imaging, optical spectroscopy, and {\\it ROSAT} HRI observations of the giant FR II radio galaxy NVSS 2146+82.  This galaxy, which was discovered by the NRAO VLA Sky Survey (NVSS), has an angular extent of nearly 20\\arcmin\\ from lobe to lobe.  The radio structure is normal for an FR~II source except for its large size and regions in the lobes with unusually flat radio spectra.  Our spectroscopy indicates that the optical counterpart of the radio core is at a redshift of $z=0.145$, so the linear size of the radio structure is $\\sim$4$h_{50}^{-1}$ Mpc, $H_{0} = 50h_{50}$. This object is therefore the second largest FR II known (3C 236 is $\\sim$6$h_{50}^{-1}$ Mpc).  Optical imaging of the field surrounding the host galaxy reveals an excess number of candidate galaxy cluster members above the number typically found in the field surrounding a giant radio galaxy.  WIYN HYDRA spectra of a sample of the candidate cluster members reveal that six share the same redshift as NVSS 2146+82, indicating the presence of at least a ``rich group'' containing the FR II host galaxy.  {\\it ROSAT} HRI observations of NVSS 2146+82 place upper limits on the X-ray flux of $1.33\\times10^{-13}$ ergs cm$^{-2}$ s$^{-1}$ for any hot IGM and $3.52\\times10^{-14}$ ergs cm$^{-2}$ s$^{-1}$ for an X-ray AGN, thereby limiting any X-ray emission at the distance of the radio galaxy to that typical of a poor group or weak AGN.  Several other giant radio galaxies have been found in regions with overdensities of nearby galaxies, and a separate study has shown that groups containing FR~IIs are underluminous in X-rays compared to groups without radio sources.  We speculate that the presence of the host galaxy in an optically rich group of galaxies that is underluminous in X-rays may be related to the giant radio galaxy phenomenon. ",
        "introduction": "The ``giant'' radio galaxies (GRGs), which we define as double radio sources whose overall projected linear extents exceed 2$h_{50}^{-1}$ Mpc, are interesting as extreme examples of radio source development and evolution. Members of this class, which comprise only a few percent of all powerful extragalactic radio sources, have been documented for almost 25 years (e.g., \\cite{wil74}).  They have been used to constrain the spectral aging and evolution of radio sources and as tests for the evolution of conditions in intergalactic environments on Mpc scales (\\cite{str80}; \\cite{sub93}; \\cite{cot96}).  Their 1.4 GHz radio powers are generally in the regime $10^{24.5}< P_{1.4}<10^{26}$ $h_{50}^{-2}$ W Hz$^{-1}$, just above the transition between Fanaroff-Riley Types I (plumed) and II (lobed) radio structures (\\cite{fan74}).  It is unclear whether the giant sources are examples of unusually long-lived (and directionally stable) nuclear activity in radio-loud systems, or of the development of sources in unusually low-density environments.  Because of their large angular sizes, nearby giant radio galaxies can in principle be studied in great detail, but their largest-scale structures may be over-resolved and undersampled by interferometers.  They have traditionally been discovered through sky surveys with compact interferometers or single dishes at relatively low frequencies, where angular resolution is modest but large fields of view and diffuse steep-spectrum structures can be imaged more easily.  The source NVSS 2146+82 was noted as a candidate giant radio galaxy when it appeared in the first 4$^\\circ$\\ by 4$^\\circ$\\ field surveyed by the NRAO VLA Sky Survey (NVSS: \\cite{Condon98}), a northern-hemisphere survey at 1.4 GHz using the VLA D configuration at 45\\arcsec\\ (FWHM) resolution.  Figure \\ref{fig:NVSScontours} shows contours of the NVSS image at 45\\arcsec\\ resolution.  There are two symmetric, extended lobes (D and E) on either side of an unresolved component C, plus an unusually large number of other radio sources within 10\\arcmin\\ of C.  Two of these (A and B) are also symmetrically located around C.  Comparison with the Digital Sky Survey (DSS) showed that source C coincides with an $\\sim$18$^{\\rm th}$ mag elliptical galaxy to within the uncertainties in the NVSS and DSS positions.  If the elliptical galaxy is the host of an unusually large radio source (C+D+E), then the apparent magnitude suggests that the whole structure may be similar in linear scale to 3C\\,236.  The DSS also shows a nearby image that might be another galactic nucleus, and a faint extended feature suggesting a possible ``tail'' or interaction.  We have undertaken several observational studies of the radio and optical objects in the field to determine their nature and to clarify the relationships between the optical and radio sources.  These studies include:  \\begin{enumerate} \\item High resolution radio imaging at 4.9 and 8.4 GHz to locate any compact flat-spectrum radio components in the field, and thus to identify any AGN that could be responsible for some or all of the other radio emission, \\item A search for fainter diffuse radio emission between the D and E components that might link them together or to other sources in the field and thus clarify their physical relationship, \\item Higher-resolution radio imaging of the other radio sources in the field to explore whether they might be physically related to the diffuse components, or to each other by gravitational lensing, \\item Optical spectroscopy of both optical ``nuclei'' and other galaxies in the field, \\item UBVRI optical photometry of the field, and \\item X-ray imaging using {\\it ROSAT} HRI observations to search for any hot X-ray emitting gas which might be associated with an overdensity of galaxies or non-thermal X-ray emission from an AGN. \\end{enumerate}  Throughout this paper, we assume a Hubble constant $H_{0} = 50h_{50}$ km s$^{-1}$ Mpc$^{-1}$ . At a redshift of $z = 0.145$, the angular diameter distance to the radio galaxy is 708.4$h_{50}^{-1}$ Mpc, the luminosity distance is 928.7$h_{50}^{-1}$ Mpc, and $1\\arcmin$ corresponds to 206$h_{50}^{-1}$ kpc. ",
        "conclusions": "We have presented multi-wavelength observations of the unusually large FR II radio galaxy NVSS 2146+82.  The overall size of the radio source is $4h_{50}^{-1}$ Mpc, making it the second largest known FR~II source.  We have found the host galaxy to be similar in both luminosity and morphology to a sample of other low redshift FR II galaxies.  Emission line profiles seen in the spectrum of the host galaxy are double peaked, which may indicate that the ionized gas may be being accelerated by the bipolar radio jet.  We have also found evidence for an anomalously rich group of galaxies at the same redshift as NVSS 2146+82 that has little associated X-ray emitting gas.  Though unusual in having a rich environment, this source is similar to other low redshift FR IIs in clusters; the NVSS 2146+82 group is underluminous in X-rays compared to clusters of similar richness that contain no FR II.  The large radio size, lack of significant Faraday rotation and non detection of X-rays all suggest that in spite of the richness of the cluster in which this galaxy resides, it has a low gas density.  There is some morphological evidence that the host galaxy of NVSS 2146+82 may be undergoing tidal interaction with one or more of its nearest companions.  Also, an interaction may be responsible for the double-peaked emission line profiles, however the spatial resolution of the spectrum of the nucleus is not high enough to distinguish between a merger origin or radio jet/cloud interaction origin for the peculiar profiles.  Apart from the radio spectral index anomaly, the radio properties of this source are like a normal FR II source scaled up by a factor of ten, preserving the standard overall morphology and polarization structure.  In the outer regions of the source the magnetic field is likely to be so weak that inverse Compton losses to the Cosmic Microwave Background dominate synchrotron losses."
    },
    "0002/astro-ph0002205_arXiv.txt": {
        "abstract": "We present the first part of the first large and homogeneous CCD color-magnitude diagram (CMD) data base, comprising 52 nearby Galactic globular clusters (GGC) imaged in the $V$ and $I$ bands using only two telescopes (one for each hemisphere). The observed clusters represent $75\\%$ of the known Galactic globulars with $(m-M)_V\\leq 16.15$~mag, cover most of the globular cluster metallicity range ($-2.2 \\leq {\\rm [Fe/H]}\\leq -0.4$), and span Galactocentric distances from $\\sim1.2$ to $\\sim18.5$ kpc.  In this paper, the CMDs for the 39 GGCs observed in the southern hemisphere are presented. The remaining 13 northern hemisphere clusters of the catalog are presented in a companion paper. For four clusters (NGC~4833, NGC~5986, NGC~6543, and NGC~6638) we present for the first time a CMD from CCD data. The typical CMD span from the $22^{\\rm nd}$ $V$ magnitude to the tip of the red giant branch. Based on a large number of standard stars, the absolute photometric calibration is reliable to the $\\sim0.02$~mag level in both filters.  This catalog, because of its homogeneity, is expected to represent a useful data base for the measurement of the main absolute and relative parameters characterizing the CMD of GGCs. ",
        "introduction": "\\label{intro}  There are two main properties which make the study of the Galactic globular clusters (GGC) particularly interesting: 1) each cluster (with possible rare exceptions) is made up by a single population of stars, all born at the same time, in the same place, and out of the same material; 2) GGC stars have the oldest measurable age in the Universe, and therefore we believe they are the oldest fossil records of the formation history of our Galaxy.  Among the many tools we have to investigate the properties of a stellar population, the color-magnitude diagrams (CMD) are the most powerful ones, as they allow to recover for each individual star its evolutionary phase, giving precious information on the age of the entire stellar system, its chemical content, and its distance. This information allows us to locate the system in the space, giving a base for the distance scale, study the formation histories of the Galaxy, and test our knowledge of stellar evolution models.  In particular, the study of a large sample of simple stellar systems, as the GGCs, provides important clues to the Milky Way formation history. Recently, many studies on the relative ages of the GGCs have been presented with results at least controversial: while some authors find a notable age spread ($\\sim5$ Gyrs) among the clusters, others find that the bulk of GGCs is coeval. This controversy is surely mainly due to the heterogeneity of the data used in each study, where the combination of photographic and/or CCD data from the early epochs of solid state detectors has been frequently used. For this reason, a survey of both southern and northern GGCs has been started two years ago by means of 1-m class telescopes, i.e. the 91cm European Southern Observatory (ESO) / Dutch telescope and the 1m Isaac Newton Group (ING) / Jacobus Kapteyn telescope (JKT). We were able to collect the data for 52 of the 69 known GGCs with $(m-M)_V\\leq16.15$. Thirty-nine have been observed with the Dutch telescope (data that are presented in this paper, hereafter Paper I), and the remaining ones with the JKT (the corresponding CMDs will be presented in a companion paper, Rosenberg et al. \\cite{rosenberg00}, hereafter Paper II).  As a first exploitation of this new data base, we have conducted a GGC relative age investigation based on the best 34 CMDs of our catalog (Rosenberg et al. \\cite{rosenberg99}, hereafter Paper III), showing that most of the GGCs have the same age. We have also used our data base to obtain a photometric metallicity ranking scale (Saviane et al. \\cite{saviane00}, hereafter Paper IV), based on the red giant branch (RGB) morphology. We measured a complete set of metallicity indices, based on the morphology and position of the RGB. Using a grid of selected RGB fiducial points, we defined a function in the $(V-I)_0$, $M_{\\rm I}$, [Fe/H] space which is able to reproduce the whole set of GGC RGBs in terms of a single parameter (the metallicity). The use of this function will improve the current determinations of metallicity and distances within the Local Group.  There are many other parameters that can be measured from a homogeneous, well calibrated CMD data base: the horizontal branch (HB) level, homogeneous reddening and distances, etc. We are presently working on these problems. However, we believe it is now the time to present to the community this data base to give to anyone interested the opportunity to take advantage of it.  In the next section, we will describe the observations collected at the ESO/Dutch telescope during two runs in 1997. The data reduction and calibration is presented in Sect.~\\ref{dat}, while in Sect.~\\ref{comparison} a cross check of the calibration between the two runs is given. In order to facilitate the reader's work, we have included the main parameters characterizing our clusters in Sect.~\\ref{parameter}. Finally, the observed fields for each cluster, and the obtained CMDs are presented and briefly discussed in Sect.~\\ref{cmds}.  \\begin{figure} \\psfig{figure=h1679f1.ps,width=8cm} \\label{dist} \\caption{ Heliocentric distribution of all GGCs with $(m-M)_V\\leq16.15$~mag. In the {\\it upper panel}, the GGCs projection over the Galactic plane is presented. The {\\it open circles} represent the clusters studied in the present paper (Paper I); the {\\it open squares}, the GGCs of Paper II; and {\\it the asterisk}, the GGC Pal~1 (Rosenberg et al. \\cite{rosenberg98}). The clusters which are not included in our catalog are marked by {\\it open triangles}. In the {\\it lower panel}, the XZ projection is shown. The Milky Way is schematically represented.} \\end{figure}  \\begin{table} \\caption[]{Observed clusters at the DUTCH in 1997} \\label{list} \\begin{tabular}{rcccccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} ID & Cluster & Other & Obs. & Obs.    & Seeing & Long.\\\\ & (NGC)   & Name  & date & fields  & $V/I(\\arcsec$) & Exp.(s)\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 1 & 104  & 47 Tuc \t\t& 23/Dec & 2 &  1.4/1.3 & 1800 \\\\ 2 & 288  & -       \t \t& 24/Dec & 3 &  1.4/1.4 & 1800 \\\\ 3 & 362  & -       \t \t& 26/Dec & 3 &  1.6/1.5 & 1200 \\\\ 4 & 1261 & -        \t\t& 24/Dec & 3 &  1.3/1.3 & 1800 \\\\ \\hline 5 & 1851 & -       \t \t& 23/Dec & 2 &  1.3/1.2 & 1800 \\\\ 6 & 1904 & M 79   \t \t& 24/Dec & 3 &  1.3/1.2 & 1800 \\\\ 7 & 2298 & -        \t\t& 23/Dec & 2 &  1.3/1.2 & 1800 \\\\ 8 & 2808 & -       \t \t& 11/Apr & 2 &  1.3/1.2 & 1500 \\\\ \\hline 8 & 2808 & -       \t \t& 26/Dec & 2 &  1.5/1.4 & 1500 \\\\ 9 & -    & E3     \t \t& 23/Dec & 2 &  1.5/1.4 & 1800 \\\\ 10 & 3201 & -       \t \t& 12/Apr & 2 &  1.5/1.4 & 900 \\\\ 10 & 3201 & -       \t \t& 24/Dec & 3 &  1.3/1.2 & 900 \\\\ \\hline 11 & 4372 & -       \t \t& 13/Apr & 2 &  1.3/1.2 & 1500 \\\\ 12 & 4590 & M 68         \t& 14/Apr & 2 &  1.2/1.2 & 1500 \\\\ 13 & 4833 & -            \t& 15/Apr & 2 &  1.3/1.2 & 1500 \\\\ 14 & 5139 & $\\omega$ Cen\t& 11/Apr & 2 &  1.2/1.2 & 900 \\\\ \\hline 15 & 5897 & -\t\t\t& 12/Apr & 1 &  1.4/1.4 & 1500 \\\\ 16 & 5927 & -\t\t\t& 13/Apr & 1 &  1.3/1.2 & 1500 \\\\ 17 & 5986 & -\t\t\t& 14/Apr & 1 &  1.3/1.3 & 1500 \\\\ 18 & 6093 & M 80\t \t& 12/Apr & 2 &  1.3/1.2 & 1500 \\\\ \\hline 19 & 6101 & -\t\t\t& 15/Apr & 1 &  1.8/1.7 & 1500 \\\\ 20 & 6121 & M 4 \t \t& 13/Apr & 2 &  1.3/1.2 & 900 \\\\ 21 & 6171 & M 107\t \t& 14/Apr & 2 &  1.4/1.3 & 1500 \\\\ 22 & 6266 & M 62\t \t& 14/Apr & 1 &  1.7/1.6 & 1500 \\\\ \\hline 23 & 6304 & -\t\t\t& 15/Apr & 1 &  1.5/1.3 & 1500 \\\\ 24 & 6352 & -\t\t\t& 11/Apr & 1 &  1.4/1.3 & 1500 \\\\ 25 & 6362 & -\t\t\t& 12/Apr & 1 &  1.4/1.3 & 1200 \\\\ 26 & 6397 & -\t\t\t& 13/Apr & 2 &  1.3/1.2 & 900 \\\\ \\hline 27 & 6496 & -\t\t\t& 14/Apr & 1 &  1.4 1.2 & 1200 \\\\ 28 & 6541 & -\t\t\t& 11/Apr & 2 &  1.3/1.2 & 1200 \\\\ 29 & 6544 & -\t\t\t& 15/Apr & 1 &  1.4/1.4 & 1500 \\\\ 30 & 6624 & -\t\t\t& 12/Apr & 1 &  1.3/1.2 & 1500 \\\\ \\hline 31 & 6626 & M 28\t \t& 13/Apr & 1 &  1.2/1.1 & 1500 \\\\ 32 & 6637 & M 69\t \t& 14/Apr & 1 &  1.2/1.1 & 1200 \\\\ 33 & 6638 & -\t\t\t& 13/Apr & 1 &  1.2/1.2 & 900 \\\\ 34 & 6656 & M 22\t \t& 15/Apr & 2 &  1.2/1.2 & 1500 \\\\ \\hline 35 & 6681 & M 70 \t \t& 11/Apr & 1 &  1.3/1.2 & 1500 \\\\ 36 & 6717 & Pal 9\t \t& 12/Apr & 1 &  1.3/1.2 & 1500 \\\\ 37 & 6723 & -\t\t\t& 13/Apr & 1 &  1.2/1.1 & 1200 \\\\ 38 & 6752 & -\t\t\t& 14/Apr & 1 &  1.3/1.2 & 1200 \\\\ \\hline 39 & 6809 & M 55 \t \t& 15/Apr & 1 &  1.3/1.1 & 900 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table} ",
        "conclusions": ""
    },
    "0002/astro-ph0002496_arXiv.txt": {
        "abstract": "Gravitational potential and gravitational energy are presented in {analytical} form for homogeneous right parallelepiped. ",
        "introduction": "As it is well known, Newtonian gravitational potential, of homogeneous body with constant density $\\rho$, at point (X,Y,Z) is defined as triple integral over the body's volume: \\be \\label{UXYZ} U(X,Y,Z)=G\\rho \\int \\int \\int u(X,Y,Z,x,y,z)\\,dxdydz \\ee with \\be u(X,Y,Z,x,y,z)={[(x-X)^2+(y-Y)^2+(z-Z)^2]}^{-1/2};\\ee here $G$ stands for Newtonian constant of gravitation. In spite of almost 400-year-long attempts since Isaac Newton' times, the integral (\\ref{UXYZ}) is known in closed form (not in the series!) in quite a few cases \\cite{Dub}:\\\\ a) a piece of straight line, b) a sphere, c) an ellipsoid;\\\\ note that both cases a) and  b) may be considered as particular cases of c). As to serial solution, the integral (1) is expressed in terms of the various kinds of series for external points outside the minimal sphere containing the whole body (non-necessary homogeneous), as well as for inner points close to the origin of co-ordinates. However in this note we do not touch the problem of serial solution and are only interested in {\\it exact analytical} solution of (\\ref{UXYZ}).\\\\ Recently Kondrat'ev and Antonov \\cite{KA} have obtained the  rather sofisticate analytical formulas for the gravitational potential (and the gravitational energy) of some axial-symmetric figures,  namely homogeneous  lenses with spherical surfaces of different radii. In forthcoming paper \\cite{GREN} we present some new solutions for homogeneous bodies of revolution. Here we present  the solution of Eq. (\\ref{UXYZ}) for homogeneous right (=rectangular) parallelepiped, (RP). The great simplification available in this case is that boundaries of integration over each of variables $x,y,z$ are {\\it fixed and thus independent on other two variables} \\cite{Budak}. This allows to get results in the closed form in terms of elementary functions. To this end we widely used Mathematica \\cite{Wolf}.\\\\ Let RP with center at the origin of Cartesian co-ordinates have lengths of sides $2a,2b,2c,$ so that inside the RP, $-a\\leq x\\leq a,-b\\leq y\\leq b,-c\\leq z\\leq c$ . Then Eq. (1) may be rewritten in the following symmetrical form: \\be \\label{symm} U(X,Y,Z)=G\\rho \\int_{xm}^{xp}dx\\int_{ym}^{yp}dy\\int_{zm}^{zp}u(x,y,z)dz,\\ee \\be u(x,y,z)={(x^{2}+y^{2}+z^2)}^{-1/2};\\ee \\be \\label{bounds} xm=-a-X,xp=a-X,\\, ym=-b-Y,yp=a-Y,\\, zm=-c-Z,zp=c-Z. \\ee Note that now the boundaries (\\ref{bounds}) of integration in Eq.(\\ref{symm})  depend on $X,Y,Z,$ but this is not so \"dangerous\" as the case of dependence of boundaries on other variables of integration.\\\\ Now we are ready to start calculation of integral (\\ref{symm}) in analytical form. Certainly final result should not depend on the order of integration [3] and this may be used to check the calculation.\\\\ First, an integration over $z$ gives: \\be \\label{line} uz(X,Y,Z,x,y)=\\int_{zm}^{zp} u(x,y,z) dz=[\\ln (z+1/u)]_{z=zm}^{z=zp}. \\ee Here $u$ is the same function as in Eq. (4).\\\\ Second, an integration over $y$ gives: \\be \\label{rect} uyz(X,Y,Z,x)= \\left [ \\left[ -x\\arctan{y z u \\over x}+y\\ln (z+ 1/u)+ z \\ln(y+1/u)\\right]_{z=zm}^{z=zp}\\right ]_{y=ym}^{y=yp}.\\ee And  integration over $x$ gives the final expression for the gravitational  potential of homogeneous RP: \\be \\label{fin} U(X,Y,Z)= G\\rho \\biggl[\\biggl[\\biggl[ v(x,y,z)+v(y,z,x)+v(z,x,y) \\biggr] _{z=zm}^{z=zp}\\biggr]_{y=ym}^{y=yp}\\biggr]_{x=xm}^{x=xp};\\ee \\be v(x',y',z')=y'z' \\ln [x'+{({x'}^2+{y'}^2+{z'}^2)}^{1/2}]-{{x'}^2\\over 2}\\arctan {y'z'\\over x'{({x'}^2+{y'}^2+{z'}^2)}^{1/2}}.\\ee Here  co-ordinates $X,Y,Z$ are as given in (5). In Appendix we present the final result (8,9,5) in explicit form. \\\\ These formulas (\\ref{fin},9,5) may look cumbersome and unpractical while in fact they allow to easily (and exactly!) calculate gravitational potential of homogeneous RP at {\\it arbitrary} point {\\it inside} as well {\\it outside} the body. The payment for this universality does not seem very high. Also we should note that  obtaining the result (\\ref{fin},9,5) is non-trivial in the sense that direct using of Mathematica's command for integration of (1,2), Integrate[f[x],x], (and especially command for definite integral, Integrate[f[x],\\{x,xm,xp\\}]), gives enormous complex expressions even leaving some integrals uncalculated!\\\\ Note that a physically evident symmetry relative to   all three sign changes, $X\\rightarrow -X$, $Y\\rightarrow -Y$, and $Z\\rightarrow-Z$ may be used to check final formulas.\\\\Let us consider now the various particular cases. ",
        "conclusions": ""
    },
    "0002/astro-ph0002175_arXiv.txt": {
        "abstract": "Most of the theoretical results on the kinematic amplification of small-scale magnetic fluctuations by turbulence have been confined to the model of white-noise-like ($\\delta$-correlated in time) advecting turbulent velocity field. In this work, the statistics of the passive magnetic field in the diffusion-free regime are considered for the case when the advecting flow is finite-time correlated. A new method is developed that allows one to systematically construct the correlation-time expansion for statistical characteristics of the field such as its probability density function or the complete set of its moments. The expansion is valid provided the velocity correlation time is smaller than the characteristic growth time of the magnetic fluctuations. This expansion is carried out up to first order in the general case of a $d$-dimensional arbitrarily compressible advecting flow. The growth rates for all moments of the magnetic-field strength are derived. The effect of the first-order corrections due to the finite correlation time is to reduce these growth rates. It is shown that introducing a finite correlation time leads to the loss of the small-scale statistical universality, which was present in the limit of the $\\delta$-correlated velocity field. Namely, the shape of the velocity time-correlation profile and the large-scale spatial structure of the flow become important. The latter is a new effect, that implies, in particular, that the approximation of a locally-linear shear flow does not fully capture the effect of nonvanishing correlation time. Physical applications of this theory include the small-scale kinematic dynamo in the interstellar medium and protogalactic plasmas. ",
        "introduction": "The study of the statistics of magnetic fluctuations excited by a random Gaussian white-noise-like advecting velocity field, was pioneered by Kazantsev~\\cite{Kazantsev}, and, in more recent times, has generated a considerable amount of research (see, e.~g., Ref.~\\cite{Almighty_Chance} and references therein, as well as Refs.~\\cite{KA,Gruzinov_Cowley_Sudan,Vergassola,Rogachevskii_Kleeorin,Chertkov_etal_dynamo,BS_metric,AAS_thesis,SBK_review}). While much attention has concentrated on resistive dynamo problems, most often for very large magnetic Prandtl numbers, it is well known that the fundamental Zeldovich's, or ``stretch--twist--fold,'' mechanism of the magnetic-energy amplification (the so-called ``fast dynamo'') is active regardless of the presence of the resistive (diffusive) regularization~\\cite{Vainshtein_Zeldovich}. If the initial seed magnetic field is concentrated on the scales of the same order as the characteristic scales of the advecting velocity, the stretching and folding of the magnetic-field lines by the random flow leads to an exponential growth of the magnetic fluctuations at scales that decrease exponentially fast, until the diffusive scales are reached~\\cite{Batchelor_vort_analog,KA,AAS_thesis,SBK_review}. This scenario is common in astrophysical applications such as the turbulence in the interstellar medium or in the protogalaxy where the Prandtl number ranges from~$10^{14}$ to~$10^{22}$, giving rise to~7 to~11 decades of small (subviscous) scales available to the magnetic fluctuations~\\cite{KA,Kulsrud_etal_proto,SBK_review}. In fact, the initial diffusion-free regime may well be the only one practically important in such applications as far as the kinematic approximation is concerned, since the nonlinear saturation effects are likely to set in before the diffusion scales are reached~\\cite{Kulsrud_lecture}. On the fundamental physical level, the diffusion-free regime, in which the magnetic-field lines are fully frozen into the flow, exhibits most clearly the underlying symmetry properties of the passive advection~\\cite{BS_metric}.  With a few notable exceptions (such as Refs.~\\cite{DMSR_mean_field,MRS_dynamo_theorem,Gruzinov_Cowley_Sudan,Chertkov_etal_dynamo}), the dominant approach in the existing literature on the turbulent kinematic dynamo problem has been to study the statistics of passive magnetic fields advected by a flow $\\vxi(t,\\vx)$ whose two-time correlation function is approximated by a $\\delta$ function, $\\<\\vxi(t)\\vxi(t')\\>\\propto\\delta(t-t')$. This white-noise property of the velocity greatly simplifies matters: the evolution equations for such statistical quantities as the correlation functions and probability density function of the magnetic field can be derived in closed form and yield themselves to exact solution.  In this paper, we relax the white-noise assumption and explore the effects that arise when a finite-time correlated velocity field is introduced. This immediately raises the level of difficulty associated with solving the statistical problem. Within the theoretical framework adopted here, the difficulty can be described in the following terms. In the zero-correlation-time approximation, one essentially has to deal with only one closed differential equation that fully determines the desired statistics. Allowing for a finite velocity correlation time leads to an infinite number of interlinked integro-differential equations involving time-history integrals. These equations form an infinite open hierarchy that formally constitutes the exact description of the problem (these matters are explained in more detail in~\\ssecref{KT_method_hierarchy}). Solving this hierarchy in its entirety without additional assumptions appears to be an impossible task. The most obvious way to make progress is clearly to try a perturbative approach, i.e., to consider the kinematic dynamo problem with an advecting field whose correlation time is short, but finite. If the correlation time~$\\tcorr$ is assumed to be small, one can expect to be able to construct an expansion in the powers of~$\\tcorr$ (in what follows, we will frequently refer to it as {\\em the $\\tau$~expansion}) and calculate corrections to the growth rates of the moments of the magnetic field. This is the program that we undertake here.  We consider the one-point statistics of the passive magnetic field in the diffusion-free regime. In this context, the infinite hierarchy we have mentioned above interrelates the one-point probability density function~(PDF) of the magnetic field and an infinite set of response functionals. These are averaged multiple functional derivatives of the magnetic field with respect to the velocity field and its gradients. We develop a {\\em functional expansion method} that allows us to calculate successive terms in the $\\tau$~expansion and derive in a closed form a Fokker--Planck equation for the one-point~PDF of the magnetic field. We limit ourselves to advancing the expansion one order beyond the zero-correlation-time approximation. The result is a set of corrections to the growth rates of all moments of the magnetic field. These corrections are negative, so the growth rates are reduced.  The expansion is carried out assuming that the velocity correlation time is small and keeping the time integral of the velocity correlation function fixed. The latter constraint ensures that the dynamo growth rate remains finite when the correlation time vanishes. An alternative way, which is sometimes deemed preferable on physical grounds (see, e.g, Ref.~\\cite{Kinney_etal_2D}), is to fix the total energy of the velocity field. Since a $\\delta$-correlated velocity field must necessarily possess infinite energy, fixing the energy at a finite value leads to vanishing of the growth rates when~$\\tcorr=0$. The {\\em relative} ordering of the terms in the expansion is, however, the same, regardless of what is kept fixed, so the technical side of the expansion method is unaffected.  Our expansion technique will be given detailed treatment in the body of this paper. Here, let us rather discuss the finite-correlation-time effects that can be distilled on the basis of our approach. As it turns out, a number of new interesting phenomena manifest themselves already at the level of the short-but-finite-correlation-time approximation.  In the case of the $\\delta$-correlated advecting flow, the one-point statistics of the passive magnetic field are {\\em universal} in the sense that they only depend on one small-scale property of the velocity: the time integral of the one-point correlation tensor of its gradients, $\\int\\diff t\\<\\nabla\\vxi(t)\\nabla\\vxi(0)\\>$. The essential novelty in the case of finite correlation time is that this small-scale universality is lost on two accounts.  First, the $\\tau$~expansion exhibits a sensitive dependence on the specific shape of the time-correlation profile of the velocity field (in recent literature, this was first explicitly pointed out by Boldyrev~\\cite{Boldyrev_tcorr}; see also Refs.~\\cite{vanKampen,vanKampen_review}). Namely, multiple time integrals of products of velocity correlation functions enter the expressions for the expansion coefficients. Choosing different correlation profiles leads to order-one changes in the values of these coefficients. The root of this nonuniversality lies in the topology of the vertex-correction diagrams that contribute to the orders higher than the zeroth in the $\\tau$~expansion (see~\\ssecref{tau_discussion}).  Second, the first-order terms of the $\\tau$~expansion feature a part that arises from the {\\em fourth-order} derivatives of the velocity correlation function, i.e., from the second derivatives of the velocity field. In the one-point statistical approach, this is the first manifestation of the more general tendency that introducing finite correlation times brings into play the large-scale structure of the velocity field. A related effect is the loss of Galilean invariance due to the fact that the expansion terms also depend on the actual {\\em energy} of the velocity field, i.e., on the rms value of the sweeping velocity. Indeed, now that the trajectories of the fluid elements have a ``memory'' of themselves, which extends approximately one~$\\tcorr$ back in time, we should naturally expect that there will appear an effective ``correlation length'' of the velocity (in what regards the one-point statistics of the fields it advects) approximately equal to~$\\xi\\tcorr$. Therefore, the one-point statistics of the passive fields now depend not only on the instantaneous velocity difference between two fluid particles that meet at a given time (i.e., the velocity gradient at a point), but also on the velocity that swept them into place and on the variation of the velocity gradient over the correlation length. This appearance of first-order corrections due to the second derivatives of the flow is a new effect, which indicates, in particular, that the customary approximation used in the Batchelor regime, where the advecting velocity is assumed to be locally linear~\\cite{Batchelor,ZRMS_linear}, is only justified for the $\\delta$-correlated-in-time advecting fields.  Such are the main qualitative consequences of introducing a finite-time-correlated velocity field into the kinematic dynamo problem (or, in general, any passive-advection model). A few words are in order as to the quantitative impact of a finite correlation time on the dynamo action. As we have already mentioned, the effect of the first-order corrections is to reduce the growth rates of all moments of the magnetic field. Besides the nonuniversal dependence on the spatial and temporal structure of the velocity correlation function, the reduction depends in a universally calculable way on the usual set of parameters: the order of the moment, the dimension of space, and the degree of compressibility of the flow. The overall magnitude of this reductive effect is measured by the expansion parameter, which is of the order of~$\\tcorr\\gamma$, where $\\gamma$~is the growth rate of the magnetic energy. It is not hard to demonstrate (see~\\ssecref{growth_rates_expanded}) that~$\\tcorr\\gamma d\\sim(\\tcorr/\\teddy)^2$, where $\\teddy$~is the ``eddy-turnover'' time of the advecting turbulent velocity field and $d$~is the dimension of space. In a standard Kolmogorov-type turbulence setting, one would, of course, expect any such approximation to be valid at best marginally, since $\\tcorr \\sim \\teddy$. Astrophysical plasmas offer more variety in this respect, as their driving forces (typically supernova explosions) can, in fact, decorrelate faster than the turbulent eddies turn over~\\cite{Almighty_Chance}. In any event, the small-$\\tcorr$ expansion does not offer much more than qualitative, or, at best, semiquantitative, information about the way the dynamo action is modified by the finiteness of the correlation time. It is, of course, clear that introducing a finite correlation time cannot altogether suppress the fast-dynamo mechanism~\\cite{MRS_dynamo_theorem,Chertkov_etal_dynamo}. On the other hand, our conclusion that some reduction of the growth rate should be expected, is corroborated by numerical evidence~\\cite{Chandran_tcorr,Kinney_etal_2D,Chou} that suggests a reduction of about~40\\% to~50\\%. In fact, in~\\secref{sec_one_eddy}, we offer a semiquantitative evaluation of the finite-$\\tcorr$ correction to the growth rate which yields a reduction of approximately~40\\% in the three-dimensional case and for~$\\tcorr\\sim\\teddy$. Of course, this is at best just an indication of the well-behaved character of our expansion, rather than a truly solid quantitative confirmation of~it.  The literature on the $\\tau$~expansion and finite-correlation-time effects is not extensive. Kliatskin and Tatarskii~\\cite{Kliatskin_Tatarskii} were the first to propose the hierarchy of equations for the response functionals as a starting point for a method of successive approximations as applied to the description of waves propagating in a medium with random inhomogeneities. Vainshtein~\\cite{Vainshtein_KT} applied this method to the mean-field kinematic dynamo theory. The Kliatskin--Tatarskii method and its relation to our functional expansion method are discussed at the end of~\\ssecref{KT_method_exp}. Van~Kampen~\\cite{vanKampen} and Terwiel~\\cite{Terwiel} developed the so-called cumulant expansion method; van~Kampen's review article~\\cite{vanKampen_review} also contains a good critical survey of other $\\tau$-expansion schemes predating his work. His method was later applied in the kinematic-dynamo context by Knobloch~\\cite{Knobloch} and Chandran~\\cite{Chandran_tcorr}. Their treatment was Lagrangian and did not include any effects due to the explicit spatial dependence in the induction equation. Consequently, the nonuniversality of the $\\tau$~expansion with respect to the spatial structure of the velocity correlator was not captured. The van~Kampen method is discussed in detail in~\\ssecref{VK_method}. Parallel to our development of the functional expansion method, Boldyrev~\\cite{Boldyrev_tcorr} proposed a $\\tau$-expansion method that was based on the exact solution of the induction equation in the Lagrangian frame and offered a way to calculate the second moment of the magnetic field that elicited the nonuniversal character of the $\\tau$~expansion with respect to both temporal and spatial properties of the velocity correlation tensor. Molchanov, Ruzmaikin, and Sokoloff~\\cite{MRS_dynamo_theorem} considered the statistics of the kinematic dynamo in a renovating flow using the formalism of infinite products of random matrices. (See also~Ref.~\\cite{DMSR_mean_field} for the treatment of the kinematic mean-field dynamo in a renovating flow.) A version of their approach was later advanced by Gruzinov, Cowley, and Sudan~\\cite{Gruzinov_Cowley_Sudan}. Considerable progress was achieved in a nonperturbative way by Chertkov {\\em et al.}~\\cite{Chertkov_etal_scalar2D}, who studied the passive-scalar problem in two dimensions for arbitrary velocity correlation times. However, their method only works in the two-dimensional case.  Thus, while we now seem to have a fairly good understanding of the structure of the $\\tau$~expansion and such qualitative features as the loss of the small-scale universality, an adequate nonperturbative theory of the kinematic dynamo and passive advection in finite-time-correlated turbulent velocity fields remains an open problem.  This paper is organized in the following way. In~\\secref{sec_tau_exp}, our functional expansion method is systematically developed on the example of the simplest available passive-advection problem: that of the Lagrangian passive vector in an incompressible flow. In this model, no explicit spatial dependence is present. In~\\ssecref{KT_method_hierarchy}, \\ssecref{KT_method_Mark}, and~\\ssecref{KT_method_exp}, we present a functional formalism that allows one to systematically construct successive terms in the $\\tau$-expanded Fokker--Planck equation. The dependence of the expansion coefficients on the specific functional form of the velocity time correlation profile emerges. The expansion is carried out up to the first order in~$\\tcorr$. In~\\ssecref{VK_method}, our method is compared with the van~Kampen cumulant expansion method~\\cite{vanKampen}. We ascertain that results obtained via the van~Kampen method are consistent with ours. Finally, in~\\ssecref{tau_discussion}, we discuss the underlying structure of the $\\tau$~expansion in diagrammatic terms. In~\\secref{sec_KT_dynamo}, the general arbitrarily compressible space-dependent dynamo problem is solved with the aid of the functional expansion. At this level, the nonuniversality with respect to the spatial structure of the velocity correlations, as well as the loss of Galilean invariance, become evident. In~\\ssecref{hierarchy}, we explain the emergence of an infinite hierarchy of equations for the characteristic function and various averaged response functionals of the magnetic field in the passive dynamo problem with finite-time-correlated advecting flow. The hierarchy is advanced up to the emergence of the second-order response functions. In~\\ssecref{tau_expansion}, we construct the $\\tau$~expansion up to first order in the correlation time, which leads to a closed equation for the characteristic function of the magnetic field. In~\\ssecref{Fokker_Planck_tcorr}, we derive the Fokker--Planck equation for the one-point~PDF of the magnetic-field strength valid to first order in the correlation time. The distribution is lognormal. In~\\ssecref{growth_rates_expanded}, we calculate the rates of growth of all moments of the magnetic field with (negative) first-order corrections. Finally, in~\\secref{sec_one_eddy}, we give a semiquantitative argument that relates the expansion parameter to the ratio of the correlation and eddy-turnover times of the velocity field. We also evaluate the finite-$\\tcorr$ reduction of the magnetic-energy growth rate in a model incompressible turbulence consisting of eddies of a fixed size. In~\\Apref{ap_k_space}, we provide the basic relations that allow one to express the results we have obtained in the configuration space in terms of the spectral characteristics of the velocity field. Some of the more cumbersome technical details of the $\\tau$~expansion are exiled to~\\Apref{ap_response2}. ",
        "conclusions": "While the particular methods one employs to obtain the successive terms in the $\\tau$~expansion may vary and depend on one's taste and the specific demands of the stochastic problem at hand, the underlying structure of the $\\tau$~expansion remains the same and is rooted in the common properties of all turbulence closure problems (see, e.g.,~Ref.~\\cite{Krommes_review}). As we have stated in general terms in the introduction to this paper, and as was clear from our construction of the response-function formalism in~\\ssecref{KT_method_hierarchy}--\\ssecref{KT_method_exp} or of van~Kampen's explicit series solution~\\exref{Z_via_Z0} in~\\ssecref{VK_method}, averaged solutions of stochastic equations such as~\\eqref{model_eq_Ztilde} can be represented in terms of infinite sums of multiple time-history integrals containing products of time-correlation functions~$\\kappa(t_i-t_j)$ in the integrands. This summation can be visualized in terms of Feynman-style diagrams. The $n$-point diagrams represent the terms containing $n$~time-history integrations. As an example, \\figref{fig_diagrams} lists the three possible four-point diagrams.  It was noted by Kazantsev~\\cite{Kazantsev} (see also Ref.~\\cite{Vainshtein_Zeldovich}) that the white-noise approximation corresponds to the partial summation of all ladder-type diagrams such as the four-point one shown in \\figref{fig_diagrams}(a). The distinctive property of these diagrams is that the pairs of points $t_i$,~$t_j$ at which the time-correlation functions in the integrands of the time-history integrals are taken, can be fused without interfering with each other. No essential information is therefore lost when the time-correlation functions~$\\kappa(t_i-t_j)$ are approximated by $\\delta$~functions. However, in all orders of the $\\tau$~expansion but the zeroth, diagrams with more tangled topology appear: e.g., in the first order, these are the diagrams~\\ref{fig_diagrams}(b) and~\\ref{fig_diagrams}(c)]. Such diagrams are often referred to as the vertex corrections. Fusing points in these diagrams leads to the loss of terms that cannot be neglected~\\cite{fnote_KT_fusing_pts}. This is the context in which the emerging nonuniversality with respect to the shape of the time-correlation profile should be viewed.  In this paper, we restrict our consideration to the first-order terms in the $\\tau$~expansion. The relevant diagrams are the four-point ones shown in~\\figref{fig_diagrams}. The diagrams~\\ref{fig_diagrams}(b) and~\\ref{fig_diagrams}(c) give rise to the coefficients~$C_1$ and~$C_2$, respectively [see~\\eqref{FPEq_vKampen} and formulas~\\exref{def_C1} and~\\exref{def_C2}]. Upon changing variables $t_1\\leftrightarrow t_2$ in the diagram~\\ref{fig_diagrams}(b) and $t_1\\rightarrow t_2$, $t_2\\rightarrow t_3$, $t_3\\rightarrow t_1$ in the diagram~\\ref{fig_diagrams}(c), we see that these diagrams equally well correspond to the coefficients~$K_1$ and~$K_2$ [\\eqref{model_FPEq_1} and formulas~\\exref{def_K1} and~\\exref{def_K2}].  \\begin{figure} \\centerline{\\psfig{file=diagrams.eps}} \\vskip0.25in \\caption{\\small The three possible fourth-order diagrams. The times~$t_1$, $t_2$, $t_3$ are integration variables and may float along the axis, but their positions relative to each other and to the time~$t$ are determined by the limits of integration and therefore preserved. A dashed line connecting any two times~$t_i$ and~$t_j$ represents the time-correlation function~$\\kappa(t_i-t_j)$ in the integrand of a triple time-history integral.} \\label{fig_diagrams} \\end{figure}"
    },
    "0002/astro-ph0002343_arXiv.txt": {
        "abstract": " ",
        "introduction": "The luminosity function (LF) is one of the basic probes for studying galaxy formation. Particularly, the near-infrared $K$-band LF can be used as a tracer of the galaxy mass distribution since the near-infrared light of a quiescent galaxy is dominated by low-mass stars, and can be an approximate measure of the total stellar mass. Hierarchical structure formation models can be directly tested by observing the evolution of near-infrared LF over a significant redshift range (e.g., Kauffmann et al. 1993).  Near-infrared observations also have some practical advantages when one investigates the LF of galaxies at high redshift; it provides a relatively uniform and unbiased measure of the galaxy luminosity distribution, since the galaxy luminosity in the $K$-band is much less sensitive to the on-going star-formation activity, the scale of which may vary from galaxy to galaxy; also, the $K$-correction factor is relatively small and nearly independent of the Hubble type, even at large redshifts.  Observing the near-infrared LF in high-redshift-rich clusters may constrain the history of the evolution of the galaxy mass distribution in the highest density environment in the Universe. De Propris et al. (1999) recently investigated the evolution of the $K$-band LF of galaxies in  rich clusters at $z=$0.1--0.9. The measured LF can be fitted by the Schechter function, and the behavior of the characteristic magnitude, $K^*$, along redshift is consistent with those predicted by the passive-evolution models with a single starburst at $z=2-3$. Combining these results with the mild evolution of the mass-to-light ratio of the cluster elliptical galaxies, they concluded that the assembly of galaxies at the bright end of the LF was largely completed by $ z\\approx 1$.  Pushing these studies toward higher redshift constrains the formation history of galaxies in clusters more strongly. In a general field environment, Kauffmann and Charlot (1998) claimed a deficit of red giant elliptical galaxies at $z \\gtsim 1$. Franceschini et al. (1998) also argued that the number density of bright elliptical galaxies significantly decreases above $z \\sim 1.3$. Although there is still much debate on these subjects (e.g., Totani, Yoshii 1998), it would be interesting to also study the $K$-band LF of galaxies in rich clusters at $z \\gtsim 1$. Observing rich clusters has an advantage, because the galaxies are at a single distance and the LF can be evaluated by measuring the galaxy surface-density excess of the cluster region over the non-cluster region, while it is still a time-consuming task to spectroscopically measure the distances to each individual galaxy in the field. The disadvantage is that the uncertainty of the field correction could seriously affect the results, especially at the faintest end where the galaxy counts may be dominated by the foreground/background galaxies.  Recently, more than several clusters and cluster candidates at $z \\gtsim 1$ have been discovered (Dickinson 1995; Yamada et al. 1997; Stanford et al. 1997; Hall, Green 1998; Ben\\'\\i tez et al. 1999; Rosati et al. 1999). Most of these objects have been selected or identified by the surface density excess of the quiescent old galaxy population. Near-infrared luminosity distributions of the galaxies in these high-redshift clusters, however, have so far not been studied intensively.  In this paper, we present the $K$-band luminosity distribution of the galaxies in the clusters at $z \\sim 1.2$ near to the radio galaxy 3C 324 using the images obtained with the Subaru telescope. These clusters are recognized by Kristian et al. (1974) and by Spinrad and Djorgovski (1984), and firmly identified by Dickinson (1995). The clusters have been spectroscopically confirmed, and the surface-density excess was revealed to be due to the superposition of two systems at $z=1.15$ and $z=1.21$ (Dickinson 1997a, b).  The extended X-ray emission, whose luminosity is comparable to that of the Coma cluster, has been detected toward the direction of 3C 324 (Dickinson 1997a), which indicates that at least one of the two systems, probably the $z=1.21$ one in which 3C 324 exists, is a fairly collapsed massive system. Smail and Dickinson (1995) detected a weak shear pattern in the field that may be produced by a cluster associated with 3C 324. We describe the observation and the data reduction briefly in section 2. In section 3, we present the obtained $K$-band luminosity function and the surface number density distribution. Our conclusions and the discussions are given in the last section. ",
        "conclusions": "We presented the luminosity distribution and the surface-density distribution for the $K$-band selected galaxies in the region of the clusters at $z=1.15$ and 1.21 near to the radio galaxy 3C 324. While the bright end of the luminosity distribution can be represented by the Schechter-function-like exponential cut off with a characteristic magnitude of $K^* \\sim 18.4$ mag, a rapid decrease in the surface-density excess compared to the average field counts is also seen at the faint end below $K \\sim 20$ mag, $\\sim 1.5$ magnitude fainter than $K^*$.  Figure 9 compares the obtained value of $K^{*}$ with those of the lower-redshift clusters studied by De Propris et al. (1999). Various lines in the figure show the behavior expected for the no-evolution and the passive-evolution models with various cosmological parameters calculated by using GISSEL96 (Bruzual, Charlot 1993). Following De Propris et al., we used a 0.1 Gyr burst model with a Salpeter IMF and with solar metallicity. It can be seen that our result follows the trend of the intermediate-redshift clusters that is consistent with the passive evolution models with star-formation epoch of $z \\gtsim 2$. At $K \\ltsim$ 20 mag, the dominant population in the clusters of 3C 324 seems to be old quiescent galaxies which were formed  at least $\\sim 1$ Gyr ago from the observed epoch.  The faint-end slope of the optical and near-infrared LF of the nearby and intermediate-redshift clusters has been studied extensively (Sandage et al. 1985; Thompson, Gregory 1993; Driver et al. 1994; Kashikawa et al. 1995; Biviano et al. 1995; Secker, Harris 1996; Metcalfe et al. 1994; Barger et al. 1996; Smith et al. 1997; Wilson et al. 1997; Driver et al. 1998). Through near-infrared observations, Barger et al. (1996) give $\\alpha=-1$ for the $K$-band LF of the $z\\sim 0.31$ clusters. There is no such deficit as seen in the 3C 324 clusters at the magnitude range between $M_K^*$ and $M_K^* + 3$. Wilson et al. (1997) also give $\\alpha=-1$ and $-1.3$ for the $I$-band LF of the two clusters at $z\\sim 0.2$, respectively. Although there are some unevenness within a factor of two, no rapid drop is seen between $m_I^*$ and $m_I^*+4$.  There is evidence that the optical LF of the clusters may be bimodal and better fitted by the combination of a Gaussian distribution for the bright (giant) galaxies and the Schechter function for the faint (dwarf) galaxies rather than by the single Schechter function (Sandage et al. 1985; Thompson, Gregory 1993; Kashikawa et al. 1995; Biviano et al. 1995; Secker, Harris 1996; Metcalfe et al. 1994). Indeed, both the $B$-band and $R$-band LFs of the Coma cluster show some `gap' at $\\sim 1$ mag fainter than the peak magnitude of the bright population (Biviano et al. 1995; Secker, Harris 1996). In the $R$ band, the characteristic magnitude of the `faint' population is $\\sim 1$--$1.5$ mag fainter than the peak magnitude of the `bright' one (Secker, Harris 1996). The shape of the $H$-band luminosity function of the Coma cluster (De Propris et al. 1998) is also consistent with these results.  On the other hand, a deficit of the galaxies is seen at $\\sim$ 1.5--2.5 mag fainter than the $K^*$ in the 3C 324 clusters. Although there is still an uncertainty of $\\sim 1$ mag in the value of the $K^*$, the deficit of the galaxy seems to occur at the magnitude range where the dwarf population already begins to dominate in the Coma cluster. How can we interpret this deficit of the faint galaxies at $K \\sim 21$ in the 3C 324 clusters if it is not just a statistical fluctuation of the foreground/background galaxies and an intrinsic property of the cluster?  It may be due to luminosity segregation between the inner and outer radius from 3C 324.  The physical diameter of the ``cluster region\" (R $<$ 40'') studied in this paper is $\\sim 0.7$ Mpc, and thus should be considered to be a central part of the cluster. Driver et al. (1998) have shown that the dwarf galaxies ($M_R > M_R^*+3$) are less concentrated than the luminous galaxies in the clusters with Bautz--Morgan type III, namely, irregular less-concentrated systems. The rich clusters at high redshift may still be dynamically young and share the properties of irregular clusters seen in the local universe. If the faint galaxies in the 3C 324 clusters are more populous in the outskirt region, or distributed rather flat over the scale of a few Mpc, the excess of the surface density can be easily hidden by the numerous foreground/background galaxies.  Another possibility is the intrinsic deficiency of the faint galaxy population in the cluster(s). There is a model involving a late formation of dwarf galaxies.  Kepner et al. (1997) have shown that the UV background radiation of $J_\\nu$ = 10$^{-21}$ erg s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ at the Lyman limit at $z=3$ and evolves as $(1+z)^4$ can prevent a baryonic collapse in the dark-matter halo with a circular velocity of $\\ltsim 30$ km s$^{-1}$ by $z\\sim1.5$, although the Hubble Deep Field observation did not prove the presence of many `bursting dwarf' populations (Ferguson, Babul 1998). The circular velocity, $\\sim 30$ km s$^{-1}$, corresponds to the dynamical mass of $\\sim 10^9 M_\\odot$, or may be $L \\sim 0.01$ $L^*$ and much less luminous than the $K \\sim 21$ mag object ($\\sim 0.1$ $L^*$). In a rich cluster environment at a high redshift, however, the limiting mass could be larger, since there may be large extra contribution of UV radiation by the massive stars rapidly formed in the primordial giant elliptical galaxies, which constitute the bright end of the LF at the observed epoch, in addition to the general background field, which may be due to quasars.  The results presented in the paper are just for one region and the superposition of the two distinct systems could make the situation more complex. Clearly, it is important to extend the study of the LF to other high-redshift clusters as well as to constrain more firmly the faint-end of the LF in the intermediate-redshift clusters. \\\\ \\vspace{0.5cm}  We would like to thank the referee, Mark Dickinson for his invaluable comments. The present result is indebted to all the members of the Subaru Observatory, NAOJ, Japan. This research was supported by grants-in-aid for scientific research of the Japanese Ministry of Education, Science, Sports and Culture (08740181, 09740168). This work was also supported by the Foundation for the Promotion of Astronomy of Japan. The Image Reduction and Analysis Facility (IRAF) used in this paper is distributed by National Optical Astronomy Observatories, operated by the Association of Universities for Research in Astronomy, Inc., under contact to the National Science Foundation."
    },
    "0002/astro-ph0002318_arXiv.txt": {
        "abstract": "All 543 pointed observations of the ROSAT High Resolution Imager (HRI) with exposure times higher than 50 sec and performed between 1990 and 1998 in a field of 10\\degr\\ x 10\\degr\\ covering the Large Magellanic Cloud (LMC) were analyzed. A catalogue was produced containing 397 X-ray sources with their properties measured by the HRI. The list was cross-correlated with the ROSAT Position Sensitive Propotional Counter (PSPC) source catalogue presented by Haberl \\& Pietsch (1999) in order to obtain the hardness ratios for the X-ray sources detected by both instruments. 138 HRI sources are contained in the PSPC catalogue, 259 sources are new detections. The spatial resolution of the HRI was higher than that of the PSPC and the source position could be determined with errors  mostly smaller than 15\\arcsec\\ which are dominated by systematic attitude errors. After cross-correlating the source catalogue with the SIMBAD data base and the TYCHO catalogue 94 HRI sources were identified with known objects based on their positional coincidence and X-ray properties. Whenever more accurate coordinates were given in catalogues or literature for identified sources, the X-ray coordinates were corrected and the systematic error of the X-ray position was reduced. For other sources observed simultaneously with an identified source the coordinates were improved as well. In total the X-ray position of 254 sources could be newly determined. The catalogue contains 39 foreground stars, 24 supernova remnants (SNRs), five supersoft sources (SSSs), nine X-ray binaries (XBs), and nine AGN well known from literature. Another eight sources were identified with known candidates for these source classes. Additional 21 HRI sources are suggested in the present work as candidates for SNR, X-ray binary in the LMC, or background AGN because of their extent, hardness ratios, X-ray to optical flux ratio, or flux variability. ",
        "introduction": "The Magellanic Clouds (MCs) as the nearest galaxies to the Milky Way allow us to resolve their stellar content in various wavelength bands. X-ray observations combined with optical and radio data can be used to investigate the physical properties of individual X-ray sources as well as the statistical properties of different source classes in a galaxy as a whole. The quantitative and positional distribution of X-ray sources in the MCs will help us to understand the unresolved X-ray emission from more distant galaxies.  After the first observation of X-ray emission from the MCs in 1968 (Mark et al.\\ 1969) four permanent (LMC X-1, X-2, X-3, and X-4, Leong et al.\\ 1971; Giacconi et al.\\ 1972) and few transient X-ray sources were found in the LMC in several satellite missions (UHURU, SAS-3, Copernicus, Ariel-V, HEAO-1). An extensive pointed survey of the LMC was performed by the Einstein Observatory between 1979 and 1981. The two detectors on board this satellite, the Imaging Proportional Counter and the High Resolution Imager, were sensitive enough to detect X-ray binaries, SSSs, and SNRs at the distance of the LMC (55kpc). Long et al.\\ (1981) published a list of 97 discrete X-ray sources in the direction of the LMC and the same data was re-analyzed by Wang et al.\\ (1991) finally giving a list of 105 sources. 54 discrete X-ray sources were identified with objects in the LMC, most of the remaining sources were associated with foreground stars and background AGN. In EXOSAT observations few additional X-ray sources were found (Jones et al.\\ 1985; Pakull et al.\\ 1985; Pietsch et al.\\ 1989).  The next thorough survey of the LMC was made by ROSAT in the energy range of 0.1 -- 2.4 keV (Tr\\\"umper 1982). From 1990 to 1998 ROSAT performed more than 700 pointed observations in a 10 by 10 degree field centered on the LMC. Haberl \\& Pietsch (1999b, hereafter HP99b) analyzed 212 PSPC observations and created a catalogue of 758 X-ray sources.  In this work results of the analysis of the ROSAT HRI data of the LMC are presented. A description of the HRI detector can be found in David \\et\\ (1996). A source catalogue was obtained in a similar way as in HP99b and many sources were identified by cross-correlating the source list with other existing catalogues. With the help of known properties of different source classes we looked for new candidates for SNRs, stars, and hard X-ray sources which mainly consist of X-ray binaries and absorbed background AGN. ",
        "conclusions": "The analysis of all 543 ROSAT HRI pointed observations performed between 1990 and 1998 with exposure times higher than 50 sec is presented. Using a maximum likelihood algorithm the source detection resulted in a catalogue of 397 sources which was cross-correlated with the SIMBAD data base and the TYCHO catalogue. Further X-ray properties could be obtained for HRI sources contained in the PSPC catalogue of HP99b.  The high spatial resolution of the HRI enabled the identification of 94 HRI sources with well known objects based on the positional coincidence and considering their extent and hardness ratios. The coordinates of most of the identified sources could be improved to more accurate positions and allowed the positional correction of other HRI sources. Thus for 254 sources the systematic error for their coordinates could be reduced to values smaller than 7\\arcsec\\ which is the standard systematic error of ROSAT observations.  For different source classes like SSS, X-ray binary, SNR, Galactic stars, and background AGN classification criteria could be derived from the extent and hardness ratios of the identified sources. We looked for flux variability of the sources and for likely optical counterparts. Five newly detected HRI sources were classified as candidates for SNRs because of their extent, two HRI sources which were identified with an LMC O and a B star as HMXB candidates. Eleven sources with probable bright optical counterpart and small X-ray to optical flux ratio are classified as stellar sources. Three sources with hard and/or absorbed X-ray spectrum indicated by the PSPC hardness ratios are likely candidates for X-ray binaries or AGN. Two of the hard X-ray sources show flux variability and for each of these an optical counterpart was found.  With the help of HRI observations many new X-ray sources were found. Further follow-up observations in X-ray, optical, or radio wavelengths with spectral information are needed to characterize these sources in more detail."
    },
    "0002/astro-ph0002404_arXiv.txt": {
        "abstract": "A sample of 37 nearby galaxies displaying composite LINER/H~II and pure H~II spectra was observed with the VLA in an investigation of the nature of their weak radio emission.  The resulting radio contour maps overlaid on optical galaxy images are presented here, together with an extensive literature list and discussion of the individual galaxies. Radio morphological data permit assessment of the ``classical AGN'' contribution to the global activity observed in these ``transition'' LINER galaxies. One in five of the latter objects display clear AGN characteristics: these occur exclusively in bulge-dominated hosts. ",
        "introduction": "LINER galaxies are a class of active galaxies characterized by the presence of strong, nuclear, low-ionization emission lines (Heckman 1980).  Like Seyfert nuclei, LINERs emit much stronger optical low-ionization forbidden lines compared to H~II nuclei, whose line emission is powered by photoionization from young massive stars, but LINERs have a characteristically lower ionization state than Seyferts.  In a recent optical survey of nearly 500 nearby galaxies, Ho, Filippenko, \\& Sargent (1995, 1997a, 1997b) find that $\\sim$20\\% of all galaxies brighter than $B_T$ = 12.5\\,mag display LINER-type spectra.  An additional 13\\% of the objects show spectra intermediate between those of ``pure'' LINERs and ``pure'' H\\,II nuclei. Ho, Filippenko, \\& Sargent (1993; see also Ho 1996) dubbed this class as ``transition objects,'' and they hypothesized that these might be composite systems in which the optical signal of a weak active nucleus (the LINER component) has been spatially blended by circumnuclear star-forming regions (the H~II component).  Models have shown that photoionization by hot stars are able to reproduce the optical spectra of LINERs (Terlevich \\& Melnick 1985), especially those of the transition-type variety (Shields 1992; Filippenko \\& Terlevich 1992). Other researchers have advocated shock-wave heating as a viable excitation mechanism for LINERs (Fosbury \\etal 1978; Heckman 1980; Dopita \\& Sutherland 1995).  At the same time, an increasing body of evidence suggests that a significant fraction of LINERs do contain irrefutable AGN characteristics (see reviews by Filippenko 1996 and Ho 1999b).  Radio surveys, in particular, have shown that many LINERs exhibit weak nuclear radio emission (Heckman 1980; Sadler, Jenkins, \\& Kotanyi 1989; Wrobel \\& Heeschen 1991; Slee \\etal 1994; Falcke, Wilson, \\& Ho 1997; Van Dyk \\& Ho 1998).  As discussed in detail by Ho (1999a), weak nuclear radio emission in early-type (elliptical and S0) galaxies is most likely related to accretion-driven energy production in these objects, but at a very low level compared to powerful radio galaxies, which are also usually hosted by early-type systems.  Radio observations provide an expecially attractive tool to further our understanding of the LINER phenomenon.  Since radio frequencies are not plagued by dust obscuration and photoelectric absorption, high-resolution radio observations potentially offer the cleanest and most efficient method to isolate an AGN core.  Although radio emission also arises from supernova remnants and H\\,II regions, AGN cores can be identified through their high degree of central concentration, small angular size, flat spectral indices, and high brightness temperatures.  Given that LINERs are so numerous, they may be the most common type of AGN known.  Their large space densities could make a major impact on the faint end of the local AGN luminosity function, which in turn has ramifications for many astrophysical issues ranging from the cosmological evolution of AGN to the contribution of AGN to the cosmic X-ray background.  We are therefore conducting high-resolution radio observations of selected LINER samples in order to investigate the interrelation between LINERs and classical AGN.  For the present paper --- the first in a series --- we have concentrated on transition-type LINERs.  We report on sensitive Very Large Array (VLA)\\footnote{The VLA is a facility of the National Radio Astronomy Observatory (NRAO) which is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} C-array observations of a sample of 25 such transition LINERs at 8.4\\,GHz.  The resulting images, having angular resolution of about 2\\arcsecpoint5, allow us to assess the strength and spatial distribution of the nuclear radio emission.  For comparison, we have also observed a small sample of 12 pure H\\,II nuclei.  In the following discussion and in upcoming papers, we adopt a Hubble constant of H$_{0}$ = 75~km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "A sample of composite LINER/H\\,II galaxies and pure H\\,II nuclei was studied using the VLA (C-array) at 8.4\\,GHz. On the basis of their radio morphological properties, the composite sources can be divided into objects with low surface brightness emission confined mainly to the plane of the galaxy and objects displaying compact nuclear components. The former objects are similar in radio morphology to the H\\,II nuclei in our sample, whereas the latter show morphologies consistent with AGN. All five of the LINER AGN are hosted by bulge-dominated galaxies (Hubble type E--Sb), and four of them show flat spectral indices between 8.4 and 1.4\\,GHz. In terms of radio luminosity, the present LINER AGN populate the range of low core radio luminosities, which implies that classical AGN of low luminosity exist in a wide range of galaxy types."
    },
    "0002/astro-ph0002297_arXiv.txt": {
        "abstract": "Since radio sources with Ultra Steep Spectra (USS; $\\alpha \\lesssim -1.30; S \\propto \\nu^{\\alpha}$) are efficient tracers of high redshift radio galaxies (HzRGs), we have defined three samples of such USS sources using the recently completed WENSS, TEXAS, MRC, NVSS and PMN radio-surveys. Our combined sample contains 669 sources with $S_{1400} > 10$~mJy and covers virtually the entire sky outside the Galactic plane ($|b|>15$\\arcdeg). For our 2 largest samples, covering $\\delta > -35\\arcdeg$, we selected only sources with angular sizes $\\Theta < 1\\arcmin$. For 410 sources, we present radio-maps with 0\\farcs3 to $\\sim$5\\arcsec\\ resolution from VLA and ATCA observations or from the FIRST survey, which allows the optical identification of these radio sources.  Comparison with spectrally unbiased samples at similar flux density levels, shows that our spectral index, flux density, and angular size selections do not affect the angular size distribution of the sample, but do avoid significant contributions by faint foreground spiral galaxies. We find that the spectral index distribution of 143,000 sources from the WENSS and NVSS consists of a steep spectrum galaxy and a flat spectrum quasar population, with the relative contribution of flat spectrum sources doubling from $S_{1400}>0.1$~Jy to $S_{1400}>2.5$~Jy. The identification fraction of our USS sources on the POSS ($R \\lesssim 20$) is as low as 15\\%, independent of spectral index $\\alpha < -1.30$. We further show that 85\\% of the USS sources that can be identified with an X-ray source are probably contained in galaxy clusters, and that $\\alpha < -1.6$ sources are excellent Galactic pulsar candidates, because the percentage of these sources is four times higher in the Galactic plane.  Our sample has been constructed to start an intensive campaign to obtain a large sample of high redshift objects ($z>3$) that is selected in a way that does not suffer from dust extinction or any other optical bias. ",
        "introduction": "Radio galaxies have now been found out to redshifts of $z=5.19$ (\\cite{wvb99b}) and radio-loud quasars out to $z=4.72$ (\\cite{hoo98}). Although new optical selection techniques such as color-dropouts, deep spectroscopy of blank fields, and narrow-band \\Lya\\ imaging have now found galaxies at similar (\\cite{ste99}) and even higher redshifts (up to $z \\sim 5.75$; \\cite{dey98}; \\cite{wey98}; \\cite{spi98}, \\cite{hu99}), radio sources are still the only objects that can be selected uniformly over all redshift ranges, and in a way that does not suffer from optical biases such as dust extinction, which is known to be important at these high redshifts (\\eg\\ \\cite{hug98}; \\cite{ivi98}; \\cite{dic98}).  At low to moderate redshift ($z \\lesssim 1$), powerful radio sources are uniquely identified with massive ellipticals (\\cite{lil84}; \\cite{owe89}; \\cite{bes98}; \\cite{mcl00}). The strongest indications that this is also true at higher redshifts comes from the near-IR Hubble $K-z$ diagram of radio galaxies which shows a remarkably close correlation from the present out to $z=5.19$ (\\cite{lil89}; \\cite{eal97}; \\cite{wvb98}, \\cite{wvb99b}). This suggests that we can use radio galaxies to study the formation and evolution of the most massive galaxies, which, by their implied star-formation history, can put important constraints on galaxy formation models, and even on cosmological parameters (\\eg \\cite{dun96}; \\cite{spi97}). Although the unification model for radio galaxies and quasars (\\eg \\cite{bar89}) suggests we could also use quasars as tracers, a detailed stellar population study of quasar host galaxies is almost impossible due to the extreme luminosity of the AGN. Furthermore, samples of radio sources designed to find large quantities of quasars require additional optical selections (\\eg\\ \\cite{gre96}, \\cite{hoo98}, \\cite{whi00}).  \\begin{figure} \\psfig{file=ds1811f1.eps,width=8.8cm} \\caption[zalpha.ps]{$\\alpha_{1400}^{325}$ against $z$ for 2 samples without spectral index selection (3CR, \\cite{spi85} and MRC, \\cite{mcc96}), and 2 USS samples (4C, \\cite{cha96a} and our new WN/TN samples, as defined in this paper). Note that the correlation is present in the spectrally unbiased 3CR and MRC, and that the 4C and our new USS samples are finding three to five times more $z>2$ radio galaxies than the MRC. The horizontal dotted line indicates the $\\alpha_{1400}^{325} < -1.3$ cutoff used in our USS sample. \\label{zalpha}} \\end{figure}  Considerable effort has been spent over the last decade to find these high redshift radio galaxies (HzRGs), which has lead to the discovery of more than 140 radio galaxies at redshifts $z>2$ (see \\eg \\cite{deb98a} for a recent summary). However by $z>3$, their numbers become increasingly sparse, and using flux limited radio surveys such as the 3CR ($S_{178} > 10$~Jy; \\cite{lai83}), or the MRC strip ($S_{408} > 0.95$~Jy; \\cite{mcc96}), the highest redshift radio galaxy found so far is at $z\\sim 3.2$ (Fig. \\ref{zalpha}; \\cite{raw90}; \\cite{mcc96}). This redshift limit arises because radio power is correlated with redshift in bright flux limited samples, and an upper limit exists in the radio luminosity. Lowering the flux limit would not only substantially increase the number of sources in these samples, but at the same time the fraction of luminous very high redshift radio galaxies would decrease (\\cite{blu98}, \\cite{jar99}). This fractional decrease would arise even if there is no decrease in co-moving space density at $z \\sim 2.5$. Such a redshift cutoff has been suggested by Bremer \\etal (1998), but recently Jarvis \\etal (1999) rule out a break at $z\\lesssim 2.5$. To efficiently find large numbers of HzRGs in acceptable observing times, it is therefore necessary to apply additional selection criteria, at the expense of completeness.  By far the most successful selection criterion has been the ultra steep spectrum criterion (\\eg \\cite{rot94}; \\cite{cha96a}; \\cite{blu98}). Selecting sources with very steep radio spectra increases dramatically the chance of finding $z>2$ radio galaxies (Fig. \\ref{zalpha}). This technique is based on the results of Tielens \\etal\\ (1979) \\nocite{tie79} and Blumenthal \\& Miley (1979)\\nocite{blu79}, who found that the identification fraction on the POSS ($R \\lesssim 20$) decreases with steepening spectral index, consistent with the steeper sources being at higher redshifts. It is now getting clear that this correlation can be explained by a combination of a K-correction of a concave radio spectrum and an increasing spectral curvature with redshift (\\cite{kro91}, \\cite{car99}; \\cite{wvb99a}). To further investigate the $z - \\alpha$ correlation, we have calculated  spectral indices using the flux densities from the WENSS (\\cite{ren97}) and NVSS (\\cite{con98}) catalogs for four different samples: the flux density limited 3CR (\\cite{spi85}) and MRC (\\cite{mcc96}) surveys, and the USS samples from the 4C (\\cite{cha96a}) and the one presented in this paper. The results (Fig. \\ref{zalpha}) show a trend for steeper spectral index sources to have higher redshifts in flux limited, spectrally unbiased samples, confirming the empirical relation out to the highest redshifts. The efficiency of the USS criterion is clearly illustrated by the fact that the 4C USS sample (\\cite{cha96a}) contains 50\\% $z>2$ sources, and by the early spectroscopic results on the USS samples presented in this paper, which indicate that $\\sim$2/3 of our sources have $z>2$. It is even more impressive to note that 13 of the 14 radio galaxies at $z>3.5$ we know of have been found from samples with a steep spectral index selection\\footnote{The only exception is VLA J123642+6213 (\\cite{wad99}), which has been identified in the HDF, but it does have a steep spectral index ($\\alpha_{1400}^{8500}=-0.94$).}! The limitation of this technique is that the steepest spectrum sources are rare, comprising typically only 0.5\\% (at $\\alpha < -1.30$) of a complete low frequency sample; therefore, large and deep all sky surveys are needed to obtain a significant sample of USS sources.  With the advent of several new deep all-sky surveys (\\S 2), it is now possible for the first time to construct a well defined all-sky USS sample with optimized selection criteria to find large numbers of $z>3$ radio galaxies. In this paper, we describe the construction of such a sample, and present high resolution radio observations needed to determine accurate positions and morphologies. This information is essential for the optical and near-IR identifications, and subsequent optical spectroscopy of a significant sub-set of our sample, which will be described in a future papers. The organization of the paper is as follows: we describe the radio surveys we used in \\S 2 and define our samples in \\S 3. We present and discuss our radio observations in \\S 4, and present our conclusions in \\S 5. \\begin{table*} \\centerline{\\bf Table 1: Radio Surveys}  \\footnotesize \\begin{tabular}{lccc} \\hline \\\\ & WENSS & TEXAS & MRC \\\\ \\\\ \\hline \\\\ Frequency (MHz) & 325 & 365 & 408 \\\\ Sky region (J2000) & $\\delta >$ +29\\arcdeg & $-$35\\fdg7 $< \\delta <$ +71\\fdg5 & $-$85\\arcdeg $< \\delta <$ +18\\fdg5 \\\\ \\# of sources & 229,576 & 67,551 & 12,141 \\\\ Resolution & $ 54\\arcsec \\times$ 54\\arcsec cosec$\\delta$ & 10\\arcsec$^a$ & $2\\farcm62 \\times 2\\farcm86 \\sec (\\delta - 35\\fdg5)$ \\\\ Position uncertainty & 1\\farcs5 & 0\\farcs5---1\\arcsec & 8\\arcsec \\\\ (strong sources) & & & \\\\ RMS noise & $\\sim$4 mJy & 20 mJy & 70 mJy \\\\ Flux density limit & 18 mJy & 150 mJy & 670 mJy \\\\ Reference & \\cite{ren97} & \\cite{dou96} & \\cite{lar81} \\\\ \\\\ \\hline \\hline \\\\ & NVSS & FIRST & PMN \\\\ \\\\ \\hline \\\\ Frequency (MHz) & 1400 & 1400 & 4850 \\\\ Sky region (J2000) & $\\delta > -$40\\arcdeg$^b$ & $7^h20^m < \\alpha < 17^h20^m$; +22\\fdg2 $< \\delta <$ +57\\fdg5 & $-$87\\fdg5 $< \\delta <$ +10\\arcdeg \\\\ & & $21^h20^m < \\alpha < 3^h20^m$; $-$2\\fdg5 $< \\delta <$ +1\\fdg6 \\\\ \\# of sources & 1,689,515 & 437,429 & 50,814 \\\\ Resolution & 45\\arcsec $\\times$ 45\\arcsec & 5\\arcsec $\\times$ 5\\arcsec & 4\\minpoint2 \\\\ Position uncertainty & 1\\arcsec & 0\\farcs1 & $\\sim$45\\arcsec \\\\ (strong sources) & & & \\\\ RMS noise & 0.5 mJy & 0.15 mJy & $\\sim$8 mJy \\\\ Flux density limit & 2.5 mJy & 1 mJy & 20 mJy \\\\ Reference & \\cite{con98} & \\cite{bec95} & \\cite{gri93} \\\\ \\hline \\end{tabular}  $^a$ The Texas interferometer has a complicated beam. However, sources with separations between 10\\arcsec\\ and 2\\arcmin\\ can be successfully modeled as doubles, and will have a single entry in the catalog. See \\cite{dou96} for details.  $^b$  Some small gaps are not covered. They are listed on the NVSS homepage (1998 January 19 version). \\end{table*} ",
        "conclusions": "We have constructed three spatially separated samples of USS sources containing a total of 669 objects. High-resolution radio observations of more than half of these show that the median size is $\\sim$6\\arcsec, independent of 1.4 GHz flux density, which is consistent with results of similar resolution surveys of samples without spectral index selection. The absence of a downturn in angular size at the lowest fluxes indicates that we do not include significant numbers of spiral galaxies in our sample. A USS sample fainter than ours would therefore include more of these foreground sources, and be less efficient to find HzRGs.  The identification fraction on the POSS is $\\sim$15\\%, with no clear dependence on spectral index, indicating that the HzRGs in the sample are all too distant to be detected, and the POSS detections consist of different classes of objects. A correlation of our USS samples with X-ray catalogs showed that at least 85\\% of the X-ray identifications seem to be galaxy clusters known from the literature or by inspection of the POSS. We conclude that (1) the majority of the 'non HzRG' USS sources in our sample are clusters, and (2) the combined selection of USS and X-ray sources is an extremely efficient technique to select galaxy clusters.  The above results indicate that up to 85\\% of our USS sample might be HzRGs. To identify these objects, we have started an intensive program of R- and K-band imaging on 3--10m class telescopes. Initial results from optical spectroscopy indicate that 2/3 are indeed $z>2$ radio galaxies (\\cite{deb98a}), and K-band imaging of optically undetected ($R > 25$) sources (see \\eg, \\cite{wvb99a}) has already lead to the discovery of the first radio galaxy at $z>5$ (\\cite{wvb99b})."
    },
    "0002/astro-ph0002068_arXiv.txt": {
        "abstract": "In a previous paper (Yuan 1999, hereafter Paper I), we find that the outer boundary conditions(OBCs) of an optically thin accretion flow play an important role in determining the structure of the flow. Here in this paper, we further investigate  the influence of OBC on the dynamics and radiation of the accretion flow on a more detailed level. Bremsstrahlung and synchrotron radiations amplified by Comptonization are taken into account and two-temperature plasma assumption is adopted. The three OBCs we adopted are the temperatures of the electrons and ions and the specific angular momentum of the accretion flow at a certain outer boundary. We investigate the individual role of each of the three OBCs on the dynamical structure and the emergent spectrum. We find that when the general parameters such as the mass accretion rate $\\dot{M}$ and the viscous parameter $\\alpha$ are fixed, the peak flux at various bands such as radio, IR and X-ray, can differ by as large as several orders of magnitude under different OBCs in our example. Our results indicate that OBC is both dynamically and radiatively important therefore should be regarded as a new ``parameter'' in accretion disk models.  As an illustrative example, we further apply the above results to the compact radio source Sgr A$^*$ located at the center of our Galaxy. Advection-dominated accretion flow (ADAF) model has been turned out to be a great success to explain its luminosity and spectrum. However, there exists a discrepancy between the mass accretion rate favored by ADAF models in the literature and that favored by the three dimensional hydrodynamical simulation, with the former being 10-20 times smaller than the latter. By seriously considering the outer boundary condition of the accretion flow, we find that due to the low specific angular momentum of the accretion gas, the accretion in Sgr A$^*$ should belong to a new accretion pattern, which is characterized by possessing a very large sonic radius. This accretion pattern can significantly reduce the discrepancy between the mass accretion rate. We argue that the accretion occurred in some detached binary systems, the core of nearby elliptical galaxies and active galactic nuclei(AGNs) very possibly belongs to this accretion pattern. ",
        "introduction": "The dynamics of an accretion flow around a black hole is described by a set of non-linear differential equations. So this is intrinsically an initial value problem, and generally the outer boundary condition(OBC) has a significant influence on the global solution. In the standard thin disk model (Shakura \\& Sunyaev 1973), all the differential terms are neglected therefore the equations are reduced to a set of algebraic equations which don't entail any boundary condition at all. However, in many cases it has been shown that the differential terms, such as the inertial and the horizontal pressure gradient terms in the radial momentum equation and the advection term in the energy equation, play an important role and can not be neglected (Begelman 1978; Begelman \\& Meier 1982; Abramowicz et al. 1988). This is especially the case  for an optically thin advection dominated accretion flow (ADAF) (Ichimaru 1977; Rees et al. 1982; Narayan \\& Yi 1994; Abramowicz et al. 1995; Narayan, Kato \\& Honma 1997; Chen, Abramowicz \\& Lasota 1997). It is the existence of the differential terms that will make OBC an important factor to determine the behavior of an accretion flow.  This expectation has been initially confirmed in our Paper I. In that paper, taking optically thin one- and two-temperature plasma as examples, we investigated the influence of OBC on the dynamics of an accretion flow. We adopted the temperature and the ratio of the radial velocity to the local sound speed (or, equivalently, the angular velocity $\\Omega_{\\rm out}$) at a certain outer boundary $r_{\\rm out}$ as the outer boundary conditions and found that in both cases, the topological structure and the profiles of angular momentum and surface density of the flow differ greatly under different OBCs. In terms of the topological structure and the profile of the angular momentum, three types of solutions are found. When $T_{\\rm out}$ is relatively low, the solution is of type I. When $T_{\\rm out}$ is relatively high and the angular velocity $\\Omega_{\\rm out}$ is high, the solution is of type II. Both types I and II possess small sonic radii, but their topological structures and angular momentum profiles are different. When $T_{\\rm out}$ is relatively high and the angular velocity is lower than a critical value, the solution is of type III, characterized by possessing a much larger sonic radius. For a one-temperature plasma or ions in a two-temperature plasma, the discrepancy of the temperature in the solutions with different OBCs lessen rapidly away from the outer boundary, but the discrepancy of the electron temperature at $r_{\\rm out}$ persists throughout the disk. This is because when the accretion rate is low, the electrons are basically adiabatic, i.e., both the local radiative cooling and the energy transfer from the ions to the electrons can be neglected compared with the energy advection of the electrons. While for a one-temperature plasma or ions, the local viscous dissipation in the energy balance plays an important role, thus their temperature is more {\\em locally} determined.  In the present paper we will focus on the two-temperature accretion flows. Our first aim is to extend our study in Paper I by including synchrotron radiation as well. This is a very strong local radiative cooling mechanism in the inner region of a disk when a magnetic field is present. We will check whether our conclusion of Paper I still holds or not in this case.  The main aim of this paper is to calculate the emergent spectrum under different OBCs and to investigate the applications. We want to probe whether and how the spectrum is dependent upon the OBCs (Sections 2 \\& 3). We find that each of the three OBCs has a significant influence on the emergent spectrum. This result indicates the importance of considering the initial physical state of the accretion flow. In Section 4 we apply our theory of OBC to Sgr A$^*$. We find that the discrepancy between the mass accretion rate favored by the numerical simulation and that required in all ADAF models in the literature is naturally solved by seriously considering the OBCs of accretion flows. The last section is devoted to summary and discussion, where the promising applications in nearby elliptical galaxies and Active Galactic Nuclei (AGNs) are shortly  discussed. ",
        "conclusions": "In this paper, we numerically solve the radiation hydrodynamic equations describing an optically thin advection dominated accretion flow and calculate its emergent spectrum. We fix all the parameters such as $\\alpha, \\beta, M$ and $\\dot{M}$, and set the outer boundary condition as $\\Ti, \\Te$ and $\\lambda_{\\rm out}(\\equiv v/c_s)$ at certain outer boundary $r_{\\rm out}$. Our primary focus is to investigate the effects of the outer boundary conditions on the structure of the global solution and, especially, on the emergent spectrum.  We find that OBCs must lie in a certain range, otherwise we can't find the corresponding global solutions. However, this range is large enough in the sense that both the structure of the solution and its corresponding spectrum differ greatly from each other under different OBCs. Three types of solutions are also found under various OBCs, qualitatively agree with the result of Paper I, where magnetic field is neglected and the description to radiation is crude. The value of the sonic radius in type III solution is much larger than types I and II.  We examine the individual influence of $\\Ti, \\Te$ and $\\lambda_{\\rm out}$ on the structures and spectrum and find that each factor especially the temperature, plays a significant role. The peak flux in the radio, IR and X-ray  bands can differ by nearly two orders of magnitude or more in our example, although all the other parameters are exactly the same.  We also investigate the effect of modifying the value of $r_{\\rm out}$. We find that the feasible ranges of $\\Ti$ and $\\Te$ lessen with the increasing radii. If the value of $r_{\\rm out}$ is too large, the values of $\\Ti$ and $\\Te$ are almost unique. This result seems to mean the reduction of the effect of OBC on the solution. However, on the one hand we expect that in some cases, such as some binary system or the system where a thin disk-ADAF transition is expected to occur (e.g., X-ray binaries A0620-00 and V404 Cyg, and low luminosity AGN NGC 4258; see Narayan, Mahadevan, \\& Quataert 1998 and references therein), the value of $r_{\\rm out}$ might be small. On the other hand, the feasible range of $\\Omega_{\\rm out}$ is almost constant no matter how large $r_{\\rm out}$ is.  The reason why previous works on ADAF global solution (e.g., Narayan, Kato, \\& Honma 1997; Chen, Abramowicz, \\& Lasota 1997) did not find the obvious effect of OBC on the global solution is that, both Narayan, Kato \\& Honma (1997) and Chen, Abramowicz \\& Lasota (1997) concentrated on the dynamics of a {\\em one-temperature plasma}, where the {\\em local} viscous dissipation in the energy equation plays an important role. Due to this reason, their global solutions are in principle ``locally'' rather than ``globally'' determined, the effect of OBC weakens rapidly with the decreasing radii and the solutions converged rapidly over a small radial extent away from the outer boundary. We also obtained similar results with theirs for the one-temperature plasma in our Paper I. However, according to our result, the variation of $\\Omega_{\\rm out}$ across a certain critical value will produce obvious OBC-dependent behavior such as the transition of the sonic radius, no matter what type the plasma is, one- or two-temperature. They did not find this result might be due to the fact that  $\\Omega_{\\rm out}$ adopted by them, $\\Omega_{\\rm K}(r_{\\rm out})$ for Keplerian disk outer boundary condition or $\\sim 0.34 \\Omega_{\\rm K}(r_{\\rm out})$ for self-similar solution outer boundary condition, is always larger than the critical value, so the effect of OBC is very small hence is hard to find.  The present study concentrates on the low-$\\dot{M}$ case where the differential terms in the equation such as the energy advection play an important role therefore the effect of OBC are most obvious. When $\\dot{M}$ becomes higher, the role of the local radiation loss terms in the energy balance will become more important. In this case, we expect that the discrepancy due to OBC in the profile of the temperature will lessen. However, from our calculation to the one-temperature accretion flow whose temperature is also principally determined locally (Paper I), and Figure 5 in the present paper, we expect that the flow should still present OBC-dependent behavior in, e.g., the angular momentum and the Mach number profiles.  We do not include winds in the present study. Since the Bernoulli parameter of ADAF is positive therefore the gas can in principle escape to infinity with positive energy (Narayan \\& Yi 1994, 1995). Blandford \\& Begelman (1999) recently suggested that mass loss through winds might be dynamically important. The effect of winds on the spectrum of ADAF has been investigated by Quataert \\& Narayan (1999b). It is interesting to investigate the effect of OBC on ADAF with winds. We expect the result is probably similar with the present results since in that case the differential terms in the equation still play an important role.  It is a meaningful problem whether a standing shock occurs in an accretion flow. Although some authors have set up the shock-included global solutions (e.g. Chakrabarti 1996), the result is generally thought not to be so convincing because in their numerical procedure $r_{\\rm s}$ and $j$ are treated as two free parameters instead of the eigenvalues of the problem. A necessary condition for the shock formation is the existence of the global solution with a large sonic radius outside the centrifugal barrier. This is the key of the problem. According to our result, this kind of large-sonic-radius solution, belonging to our type III, can only be realized when the specific angular momentum of the accretion flow is low. We note that this requirement to the angular momentum is exactly what Narayan, Kato \\& Honma(1997) anticipated, and it has recently been confirmed by the numerical simulation by Igumenshchev, Illarionov \\& Abramowicz (1999).  The large-sonic radius solution (type III) is a new accretion pattern since in all previous studies on viscous accretion onto black holes the sonic radii are small. We find that the discrepancy in the mass accretion rate of Sgr A$^*$ between the value favored by the previous ADAF models in the literature and that favored by the hydrodynamical numerical simulation is significantly reduced if the accretion is via this pattern. This result hints us that such low angular momentum accretion may be very common in the universe. One example is the detached binary system, where the accretion material is the stellar wind from the companion therefore the angular momentum of the accreting gas is very low (Illarionov \\& Sunyaev 1975). Another example is the cores of nearby giant elliptical galaxies, where the angular momentum of the hot (accretion) gas is again assumed to be very small."
    },
    "0002/astro-ph0002012_arXiv.txt": {
        "abstract": "We introduce a new class of 2-D mass models, whose potentials are of St\\\"ackel form in elliptic coordinates. Our model galaxies have two separate strong cusps that form double nuclei. The potential and surface density distributions are locally axisymmetric near the nuclei and become {\\it highly} non-axisymmetric outside the nucleus. The surface density diverges toward the cuspy nuclei with the law $\\Sigma \\propto r^{-2}$. Our model is sustained by four general types of regular orbits: {\\it butterfly}, {\\it nucleuphilic banana}, {\\it horseshoe} and {\\it aligned loop} orbits. Horseshoes and nucleuphilic bananas support the existence of cuspy regions. Butterflies and aligned loops control the non-axisymmetric shape of outer regions. Without any need for central black holes, our distributed mass models resemble the nuclei of M31 and NGC4486B. It is also shown that the self-gravity of the stellar disc can prevent the double nucleus to collapse. ",
        "introduction": "{\\it Hubble Space Telescope} (HST) data revealed that M31 and NGC4486B have double nuclei (Lauer et al. 1996, hereafter L96; Tremaine 1995, hereafter T95). M31 has a bright nucleus (P1) displaced from the centre of the isophotal lines of outer regions and a fainter nucleus (P2) just at the centre. NGC4486B exhibits a similar structure with a minor difference: The centre of outer isophotes falls between P1 and P2. There are some explanations for the emergence of the double nuclei of these galaxies, among which the eccentric disc model of T95 has been more impressive. In the model of T95, a central black hole (BH) enforces stars to move on ``aligned\" Keplerian orbits, which may elongate in the same direction as the long-axis of the model. Stars moving on aligned Keplerian orbits linger near apoapsis and may result in P1. The mass of central ``supermassive\" BH should be much greater than the mass of neighboring disc. Otherwise, the asymmetric growth of P1 won't allow the BH to remain in equilibrium.  Goodman \\& Binney (1984) showed that central massive objects enforce the orbital structure of stellar systems to evolve towards a steady symmetric state. This result was then confirmed by the findings of Merritt \\& Quinlan (1998) and Jalali (1999, hereafter J99) in their study of elliptical galaxies with massive nuclear BHs. Within the BH sphere of influence, highly non-axisymmetric structure can only exist for a narrow range of BH mass (J99). The results of J99 show that {\\it long-axis tube} orbits of non-axisymmetric discs with central massive BHs, elongate in the both $\\pm$ directions of long-axis. Thus, the probabilities for the occurrence of two bright regions, in both sides of BH along the long-axis, are equal (these bright regions are supposed to be formed near the apogee of long-axis tubes). By this hypothesis one can interpret the double structure of NGC4486B by placing a supermassive BH between P1 and P2. However, some disadvantages arise in the case of M31. In the nucleus of M31, the formation of P1 can still be deduced from the behavior of long-axis tubes. But, there is no mathematical proof for the ``coexistence\" of P1 and P2 when the centre of P2 coincides with BH's location.  In this paper we attempt to create a model based on the self-gravity of stellar discs to show that systems with double nuclei can exist even in the absence of central BHs. Especial cases of our non-scale-free planar models are eccentric discs, which display a collection of properties expected in self-consistent cuspy systems. Our models are of St\\\"ackel form in elliptic coordinates (e.g., Binney \\& Tremaine 1987) for which the Hamilton-Jacobi equation separates and stellar orbits are regular.  In most galaxies, density diverges toward the centre in a power-law cusp. In the presence of a cusp, regular box orbits are destroyed and replaced by chaotic orbits (Gerhard \\& Binney 1985). Through a fast mixing phenomenon, stochastic orbits cause the orbital structure to become axisymmetric at least near the centre (Merritt \\& Valluri 1996). These results are confirmed by the findings of Zhao et al. (1999, hereafter Z99). Their study reveals that highly non-axisymmetric, scale-free mass models can not be constructed self-consistently. Near the cuspy nuclei, the potential functions of our distributed mass models are proportional to $r^{-1}$ as $r \\rightarrow 0$. So, we attain an axisymmetric structure near the nuclei which is consistent with the mentioned nature of density cusps. The slope of potential function changes sign as we depart from the centre and our model galaxies considerably become non-axisymmetric. Non-axisymmetric structure is supported by {\\it butterfly} and {\\it aligned loop} orbits. Close to the larger nucleus, loop orbits break down and give birth to a new family of orbits, {\\it horseshoe} orbits, which in turn generate {\\it nucleuphilic banana} orbits. Stars moving in horseshoe and banana orbits lose their kinetic energy as they approach to the nuclei and contribute a large amount of mass to form cuspy regions. ",
        "conclusions": "In his pioneering work, Euler showed the separability of motion in the potential field of two {\\it fixed} Newtonian centres of attraction. This problem was then completely solved by Jacobi (Pars 1965). It is physically impossible to keep apart these two ``point masses\", for they will attract each other leading to an eventual collapse. However, the assumed masses can be in equilibrium if they revolve around their common centre of mass (this is the classical 3-body problem). Our planar model is indeed Jacobi's problem in which we have replaced two fixed centres of gravitation with a continuous distribution of matter, where mass concentration increases towards two nuclei (P1 and P2) in power-law, strong cusps. These nuclei are maintained by an interesting family of orbits, nucleuphilic bananas. Below, we explain why the mentioned nuclei are generated and don't collapse.  The force exerted on a star is equal to $-\\nabla \\Phi$. The motion under the influence of this force can be tracked on the {\\it potential hill} of Figure 2a. This helps us to better imagine the motion trajectories.  \\subsection{The behaviour of orbits} Stars moving in nucleuphilic banana orbits undergo motions {\\it near} the 1:2 resonance. They oscillate twice in the $y$-direction for each $x$-axial oscillation. The turning points of this group of stars lie on the curves $v=v_{n,i}$. These hyperbolic curves can be chosen arbitrarily close to P1 or P2. When stars approach P1 (or P2), their motion slows down (because they climb on the cuspy portion of the potential hill and considerably lose their kinetic energy while the potential energy takes a maximum) and the orbital angular momentum switches sign somewhere on $v=v_{n,i}$. Thus, these stars spend much time in the vicinity of P1 (or P2) and deposit a large amount of mass. This phenomenon is the main reason for the generation of cuspy zones around P1 and P2. Stars moving in nucleuphilic bananas cross the $y$-axis quickly, and therefore, don't contribute much mass to the region between the nuclei.  Horseshoe orbits cause the sizes of cuspy zones to be different through the following mechanism. Stars that start their motion sufficiently close to P1 (larger nucleus), are repelled from P1 because the force vector is not directed inward in this region. As they move outward, their orbits are bent and cross the $x$-axis with non-zero angular momentum. These stars linger only near P1, and in consequence, help the cuspy zone around P1 grow more than that of P2. The asymmetry of nucleuphilic bananas, with respect to the $y$-axis, is also an origin of the different sizes of cuspy zones. In fact, horseshoe orbits are born once nucleuphilic bananas join together for $I_2=\\beta C$. Horseshoe and nucleuphilic banana orbits are the especial classes of boxlets that appropriately bend toward the nuclei. The lack of such a property in centrophobic banana orbits causes the discs of Sirdhar \\& Touma (1997) to be non-self-consistent.  Aligned loop orbits occur when the orbital angular momentum is high enough to prevent the test particle to slip down on the potential hill. The boundaries of loop orbits are defined by the ellipses $u=u_{l,1}$ and $u=u_{l,2}$. The nuclear cusps are located at the foci of these ellipses. Aligned loops have the same orientation as the surface density isocontours (compare Figures 1 and 4d). Thus, according to the results of Z99, it is possible to construct a self-consistent model using aligned loop orbits.  It is worthy to note that butterfly orbits play a significant role in maintaining the non-axisymmetric structure of the model at the moderate distances of ${\\cal O}(a)$.  \\subsection{The nature of P1 and P2} The points where the cusps have been located, are inherently unstable. With a small disturbance, stars located at ($x=\\pm a,y=0$) are repelled from these points because $-\\nabla \\Phi$ is directed outward when $r_i \\rightarrow 0$ ($i=1,2$). But, the time that stars spend near the nuclei will be much longer than that of distant regions when they move in horseshoe and banana orbits. The points P1 and P2 are unreachable, for they correspond to the energy level $E=+\\infty$. Based on the results of this paper, we conjecture that there may not be any mass concentration just at the centre of cuspy galaxies. However, a very dense region exists {\\it arbitrarily} close to the centre!  \\subsection{The double nucleus can be in equilibrium} The nuclei pull each other due to their mutual gravitational attraction and it seems that they must collapse. However, we explain that in certain circumstances, the double nucleus can be in {\\it static} equilibrium. At first we estimate the mass inside the separatrices of the surface density distribution (the mass of cuspy zones) and concentrate the matter at P1 and P2 (this is logical because the surface density distribution is almost axisymmetric near the nuclei). In this way, we obtain two point masses, $M_1$ and $M_2$. According to (\\ref{7}) and (\\ref{9}), the following relations approximately hold \\begin{eqnarray} M_i &=& \\int \\nolimits _{-\\pi}^{\\pi} \\int \\nolimits _{\\epsilon}^{r_{0i}} \\sigma _{i} r^{-1} {\\rm d}r {\\rm d}\\theta, ~\\epsilon \\rightarrow 0, \\nonumber \\\\ &=& 2 \\pi \\sigma _i \\log {\\frac {r_{0i}}{\\epsilon}}, ~~i=1,2, \\label{32} \\end{eqnarray} \\noi where $r_{0i}$ are chosen as the radii of inner tangent circles to the separatrices and the constant parameters $\\sigma _i$ are computed based on the surface density profile near the nuclei. As $\\epsilon \\rightarrow 0$, $M_i$s diverge to infinity unless a negative mechanism prevents them to grow. Consider the discs of radius $\\epsilon$ with the centres located at $(\\pm a,0)$ and call them ${\\cal D}_1$ and ${\\cal D}_2$. Since P1 and P2 are ``locally\" unstable, we rely on our previous argument that the matter is swept out from these points, allowing us to exclude ${\\cal D}_1$ and ${\\cal D}_2$ from our model for some $0< \\epsilon \\ll 1$. In this way, $M_i$s take finite values.  \\begin{figure} \\centerline{\\hbox{\\epsfxsize=1.7in\\epsfbox{fig6.ps}}} \\caption[Figure 5]{In this figure we have shown a circular ring of matter of outer regions. The masses of cuspy zones can approximately be computed and concentrated at the points ($\\pm a,0$). The gravitational attraction of the ring upon the point mass $M_1$ ($M_2$) is directed in the $x$-direction due to the existing symmetry.} \\end{figure}  $M_1$ and $M_2$ attract each other and start to move if they are not influenced by other gravitational sources. We claim that the required extra force comes from the gravitational attraction of the matter of outer regions. Consider Figure 6 where $M_1$ and $M_2$ are shown along with a ring of matter of outer regions. For brevity, we assume $\\beta=0$, which yields $M_1=M_2=M$. Due to the existing symmetry, the gravitational force exerted on $M_2$ by the assumed ring will have a resultant in the $x$-direction. When $r$ is sufficiently large, $r \\gg a$, this force is calculated as follows \\begin{eqnarray} F_x(r) &=& GM \\sigma _{\\infty} \\int\\nolimits_{-\\pi}^{\\pi} \\frac {r^{\\gamma -2} \\cos \\phi {\\rm d} \\theta} {r^2+a^2-2ar\\cos \\theta}, \\label{33} \\\\ \\cos \\phi &=& \\frac {r\\cos \\theta -a} {(r^2+a^2-2ar\\cos \\theta)^{1/2}}, \\label{34} \\end{eqnarray} where we have used $\\Sigma \\approx \\sigma _{\\infty} r^{\\gamma -3}$ with $\\sigma _{\\infty}$ being a positive constant (see Eq. (\\ref{11})). By integrating $F_x(r)$ over $r$ from some $r=R \\gg a$ to $r=\\infty$, the total force, due to the matter of outer regions, is found to be \\begin{equation} F_x = GM \\sigma _{\\infty} \\int\\nolimits_{R}^{\\infty} \\int\\nolimits_{-\\pi}^{\\pi} \\frac {r^{\\gamma -2} (r\\cos \\theta -a) {\\rm d} \\theta {\\rm d}r} {(r^2+a^2-2ar\\cos \\theta)^{3/2}}. \\label{35} \\end{equation} By a change of independent variable as $\\xi =a/r$, the integrand can be simplified in the form \\begin{eqnarray} F_x &=& b \\int\\nolimits_{0}^{\\xi _0} \\xi ^{2-\\gamma} {\\rm d}\\xi \\frac {{\\rm d}}{{\\rm d}\\xi} \\int\\nolimits_{0}^{\\pi} \\frac {{\\rm d}\\theta}{(1+\\xi ^2-2\\xi \\cos \\theta)^{1/2}}, \\label{36} \\\\ b  &=& \\frac {2GM \\sigma _{\\infty}}{a^{3-\\gamma}}, \\label{37} \\end{eqnarray} where $\\xi_0=a/R$. Consequently, \\begin{equation} F_x = b \\int\\nolimits_{0}^{\\xi _0} \\xi ^{2-\\gamma} {\\rm d}\\xi \\frac {{\\rm d}}{{\\rm d}\\xi} \\sum _{n=0}^{\\infty} \\xi ^n \\int\\nolimits_{0}^{\\pi} P_n(\\cos \\theta) {\\rm d}\\theta, \\label{38} \\end{equation} with $P_n$s being the well known Legendre functions. According to (Morse \\& Feshbach 1953) \\begin{eqnarray} \\int\\nolimits_{0}^{\\pi} P_{2k+1}(\\cos \\theta) {\\rm d} \\theta &=& 0, \\label{39} \\\\ \\int\\nolimits_{0}^{\\pi} P_{2k}(\\cos \\theta) {\\rm d} \\theta &=& \\pi \\left [ \\frac {(2k)!}{(2^k k!)^2} \\right ] ^2 \\equiv c_{2k}, \\label{40} \\end{eqnarray} one achieves \\begin{equation} F_x = b \\int\\nolimits_{0}^{\\xi _0} (\\xi ^{2-\\gamma} \\sum _{k=1}^{\\infty} 2k c_{2k} \\xi ^{2k-1}) {\\rm d}\\xi. \\label{41} \\end{equation} Integrating (\\ref{41}) over $\\xi$, yields \\begin{eqnarray} F_x &=& b Q(\\xi _0), \\label{42} \\\\ Q(\\xi _0) &=& \\xi _0 ^{2-\\gamma} \\sum _{k=1}^{\\infty} \\frac {2k c_{2k}}{2k+2-\\gamma} \\xi _0 ^{2k}. \\label{43} \\end{eqnarray} It is obvious that $Q$ is a positive function of $\\xi _0$. Therefore, from (\\ref{42}) one concludes $F_x>0$ indicating that $M_2$ is pulled away from the centre. $M_2$ will be in static equilibrium if $F_x$ is balanced with the gravitational force of $M_1$, i.e., \\begin{equation} F_x=\\frac {GM_1M_2}{(2a)^2}=\\frac {GM^2}{(2a)^2}. \\label{44} \\end{equation} By substituting from (\\ref{32}) and (\\ref{42}) into (\\ref{44}), we obtain \\begin{equation} \\delta \\equiv \\frac {\\epsilon}{r_{01}} = e^{-s},~~s>0, \\label{45} \\end{equation} where $r_{01}=r_{02}$, $\\sigma _1=\\sigma _2$ (because we assumed $\\beta=0$) and \\begin{equation} s=\\frac {4 \\sigma _{\\infty} a^{\\gamma -1} Q(\\xi _0)} {\\pi \\sigma _1}. \\label{46} \\end{equation} Our numerical computations of $\\Sigma$ reveal that $\\sigma_{\\infty} \\gg \\sigma_i$, which guarantees $\\epsilon \\ll r_{01}$ as desired. Following a similar procedure as above, one can show that the double nucleus remains in equilibrium for $\\beta >0$.  \\subsection{The nuclei of M31 and NGC4486B} In many respects, the surface density isocontours of our model galaxies are similar to the isophotal lines of the nuclei of M31 and NGC4486B. Our mass models are cuspy within two separatrices. Such curves can be distinguished in the nuclei of M31 and NGC4486B (see L96 and T95). We are not sure that the nuclei of M31 and NGC4486B are really cuspy, because existing telescopes can not highly resolve the regions around P1 and P2 (even HST images contain ``few\" bright pixels at the locations of P1 and P2). Whatever the mass distribution inside these nuclei may be, our models reveal that double nuclei can exist even in the absence of supermassive BHs.  The nucleus of M31 can also be explained by the eccentric disc corresponding to $\\beta=1$. In such a circumstance, stars moving in butterfly orbits form a local group in the vicinity of ($x=0,y=0$). The accumulation of stars around this local minimum of $\\Phi$ can create a faint nucleus like P2 (see T95). Therefore, P2 will approximately be located at the ``centre\" of loop orbits while the eccentric, brighter nucleus (P1) is at the location of the cusp. In other words, loop and high-energy butterfly orbits will control the overall shape of outer regions, horseshoe orbits will generate P1 and low-energy butterflies will create P2.  \\subsection{Challenging problems} It is not known for us if there are rotationally supported double structures or not. This idea comes from the fact that we can replace the point masses of the restricted 3-body problem (the restricted 3-body problem is usually expressed in a rotating frame) with a continuous distribution of matter. Moreover, NGC4486B and the bulge of M31, are three dimensional objects and the assumption of planar models seems to be a severe constraint.  So far we showed that the double nucleus can be in static equilibrium due to the existing gravitational effects of the model. The stability study of such states, however, remains as a challenging problem.  Our next goal is to apply the method of Schwarzschild (1979,1993) for the investigation of self-consistency."
    },
    "0002/astro-ph0002362_arXiv.txt": {
        "abstract": "We study how the internal structure of dark halos is affected if Cold Dark Matter particles are assumed to have a large cross-section for elastic collisions. We identify a cluster halo in a large cosmological N-body simulation and resimulate its formation with progressively increasing resolution. We compare the structure found in the two cases where dark matter is treated as collisionless or as a fluid. For the collisionless case the overall ellipticity of the cluster, the central density cusp and the amount of surviving substructure are all similar to those found in earlier high resolution simulations. Collisional dark matter results in a cluster which is more nearly spherical at all radii, has a {\\it steeper} central density cusp, and has less, but still substantial surviving substructure. As in the colisionless case, these results for a ``fluid'' cluster halo are expected to carry over approximately to smaller mass systems. The observed rotation curves of dwarf galaxies then argue that self-interacting dark matter can only be viable if intermediate cross-sections produce structure which does not lie between the extremes we have simulated. ",
        "introduction": "Cold dark matter scenarios within the standard inflationary universe have proved remarkably successful in fitting a wide range of observations. While structure on large scales is well reproduced by the models, the situation is more controversial in the highly nonlinear regime. Navarro, Frenk \\& White (1995, 1996, 1997; NFW) claimed that the density profiles of near-equilibrium dark halos can be approximated by a ``universal'' form with singular behaviour at small radii. Higher resolution studies have confirmed this result, finding even more concentrated dark halos than the original NFW work and showing, in addition, that CDM halos are predicted to have a very rich substructure with of order 10\\% of their mass contained in a host of small subhalos (Frenk et al 1999, Moore et al 1999a, 1999b, Ghigna et al 1999, Klypin et al 1999, Gottloeber et al 1999, White \\& Springel 1999). Except for a weak anticorrelation of concentration with mass, small and large mass halos are found to have similar structure. Many of these studies note that the predicted concentrations appear inconsistent with published data on the rotation curves of dwarf galaxies, and that the amount of substructure exceeds that seen in the halo of the Milky Way (see also Moore 1994; Flores and Primack 1994; Kravtsov et al 1998; Navarro 1998).  It is unclear whether these discrepancies reflect a fundamental problem with the Cold Dark Matter picture, or are caused by overly naive interpretation of the observations of the galaxy formation process (see Eke, Navarro \\& Frenk 1998; Navarro \\& Steinmetz 1999; van den Bosch 1999). On the assumption that an explanation should be sought in fundamental physics, Spergel \\& Steinhardt (1999) have argued that a large cross-section for elastic collisions between CDM particles may reconcile data and theory. They suggest a number of modifications of standard particle physics models which could give rise to such self-interacting dark matter, and claim that cross-sections which lead to a transition between collisional and collisionless behaviour at radii of order 10 -- 100 kpc in galaxy halos are preferred on astrophysical grounds. Ostriker (1999) argues that the massive black holes observed at the centres of many galactic spheroids may arise from the accretion of such collisional dark matter onto stellar mass seeds. Miralda-Escude (2000) argues that such dark matter will produce galaxy clusters which are rounder than observed and so can be excluded.  At early times the CDM distribution is indeed cold, so the evolution of structure is independent of the collision cross-section of the CDM particles. At late times, however, a large cross-section leads to a small mean free path and so to fluid behaviour in collapsed regions. In this Letter we explore how the structure of nonlinear objects (``dark halos'') is affected by this change. We simulate the formation of a massive halo from CDM initial conditions in two limits: purely collisionless dark matter and ``fluid'' dark matter. We do not try to simulate the the more complex intermediate case in which the mean free path is large in the outer regions of halos but small in their cores. If this intermediate case (which is the one favoured by Spergel \\& Steinhardt (1999) and by Ostriker (1999)) produces nonlinear structure intermediate between the two extremes we do treat, then our results show that collisional CDM would give poorer fits to the rotation curves of dwarf galaxies than standard collisionless CDM. Further work is needed to see if this is indeed the case. ",
        "conclusions": "An interesting question arising from our results is {\\it why} our fluid clusters have more concentrated cores than their collisionless counterparts. The density profile of an equilibrium gas sphere can be thought of as being determined by its Lagrangian specific entropy profile, i.e. by the function $m(s)$ defined to be the mass of gas with specific entropy less than $s$. The larger the mass at low specific entropy, the more concentrated the resulting profile. Thus our fluid clusters have more low entropy gas than if their profiles were similar to those of the collisionless clusters. The entropy of the gas is produced by a variety of accretion and merger shocks during the build-up of the cluster, so the strong central concentration reflects a relatively large amount of weakly shocked gas.  \\begin{inlinefigure} \\vspace*{0.3cm}\\ \\\\ \\resizebox{8cm}{!}{\\includegraphics{Figure3.eps}} \\caption{The total number of subhalos within $R_{200}$ is plotted as a function of their mass in units of $M_{200}$. Dashed and solid lines correspond to the collisionless and fluid cases respectively. Results for each simulation are plotted only for halos containing more than 40 particles. \\label{fig3}} \\end{inlinefigure}  We study gas shocking in our models by carrying out one further simulation. We take the initial conditions of S1 and replace each particle by two superposed particles, a collisionless dark matter particle containing 95\\% of the original mass and a gas particle containing 5\\%. These two then move together until SPH pressure forces are strong enough to separate them. The situation is similar to the standard 2-component model for galaxy clusters except that our chosen gas fraction is significantly smaller than observed values.  In this mixed simulation the evolution of the collisionless matter (and its final density profile) is almost identical to that in the original S1. This is, of course, a consequence of the small gas fraction we have assumed. In agreement with the simulations in Frenk et al (1999) we find that the gas density profile parallels that of the dark matter over most of the cluster but is significantly {\\it shallower} in the inner $\\sim 200 h^{-1}$kpc.  Comparing this new simulation (S1M) with its fluid counterpart (S1F) we find that in both cases the gas which ends up near the cluster centre lay at the centre of the most massive cluster progenitors at $z= 1\\sim 3$. In addition it is distributed in a similar way among the progenitors in the two cases. In Figure 4 we compare the specific entropy profiles of the cluster gas. These are scaled so that they would be identical if each gas particle had the same shock history in the two simulations. Over most of the cluster there is indeed a close correspondence, but near the centre the gas in the mixed simulation has higher entropy. (This corresponds roughly to $r < 100h^{-1}$kpc.)   \\begin{inlinefigure} \\vspace*{0.3cm}\\ \\\\ \\resizebox{8cm}{!}{\\includegraphics{Figure4.eps}} \\caption{We plot Lagrangian specific entropy profiles for the gas fluid simulation (S1F: crosses) and for the mixed simulation (S1M: open circles). In each case $m(s)$ is given in units of the individual gas particle mass, $m_g$, and the specific entropy of a particle is defined as $\\ln (m_g T_g^{1.5}/\\rho_g)$. The arrows indicate where the timescale $t_{2b}$ for 2-body heating of the gas by encounters with dark matter particles (see equation (5) of Steinmetz \\& White (1997)) is 0.1, 1, and 10 times the age of the Universe. For each $s$ we calculate $t_{2b}$ at the radius where the median specific entropy equals $s$. The dashed line with open squares is an ``entropy'' profile for S1 calculated by using the SPH kernel to calculate the density and velocity dispersion in the neighborhood of each particle, and then converting from velocity dispersion to temperature using the standard relation for a perfect monatomic gas. \\label{fig4}} \\end{inlinefigure}  As Figure 4 shows, this is partly a numerical artifact; the two entropies differ only at radii where two-body heating of the gas by the dark matter particles is predicted to be important in the mixed case. (The effect is absent in the pure fluid simulation.) The weaker shocking in the fluid case is evident from the equivalent \"entropy\" profile of S1 in Figure 4. This lies between those of the two fluid simulations, and in particular significantly above that of S1F in the central regions.   In conclusion the effective heating of gas by shocks in the fluid case is similar to but slightly weaker than that in the mixed case. This is presumably a reflection of the fact that the detailed morphology of the evolution also corresponds closely. The difference in final density profile is a consequence of three effects. In the mixed case the gas is in equilibrium within the external potential generated by the dark matter, whereas in the pure fluid case it must find a self-consistent equilibrium. In addition the core gas is heated by two-body effects in the mixed case. Finally in the pure fluid case the core gas experiences weaker shocks.  Overall our results show that in the large cross-section limit collisional dark matter is not a promising candidate for improving the agreement between the predicted structure of CDM halos and published data on galaxies and galaxy clusters. The increased concentration at halo centre will worsen the apparent conflict with dwarf galaxy rotation curves. Furthermore, clusters are predicted to be nearly spherical and galaxy halos to have similar mass in substructure to the collisionless case, although with fewer low mass subhalos. Intermediate cross-sections would lead to collisional behaviour in dense regions and collisionless behaviour in low density regions with a consequent breaking of the approximate scaling between high and low mass halos. The resulting structure may not lie between the two extremes we have simulated. Self-interacting dark matter might then help resolve the problems with halo structure in CDM models, if indeed these problems turn out to be real rather than apparent."
    },
    "0002/astro-ph0002154_arXiv.txt": {
        "abstract": "We use a 3D photoionization modeling tool to study the morpho-kinematic properties of the Planetary Nebula (PN) NGC~3132. We show that it is possible to reproduce the low resolution observations (spectra and images) with an ellipsoidal shell. However, high resolution observations, as those showing a density variation along the nebula and the \\forb{O}{3}{5007} velocity profiles, definitively rule out this description. We show that a bipolar Diabolo shape with a 40\\degre/ rotation of the symmetry axis relative to the line of sight successfully reproduces the observed images, as well as the high resolution observations. ",
        "introduction": "The determination of the three-dimensional matter distribution in planetary nebulae is essential for a knowledge of the nebula ejection mechanism. A precise understanding of the nebula geometry which produces the observed morphology can also provide constraints on stellar evolution theories. The geometry and the density distribution of PNe are generally derived from line imaging, assuming that a surface brightness enhancement corresponds to a density enhancement. After a first descriptive study of the morphology of PNe (Curtis 1918), more detailed studies indicate two different interpretations for the observed morphology of these objects: a) the different observed morphologies are due to different projections of a same common basic structure (Minkowski \\& Osterbrock 1960, Khromov \\& Kohoutek 1968); b) the different types in which PNe can be classified (bipolar, elliptical, etc.) correspond to different evolutionary stages (Balick 1987). In this second approach, the morphology can also be related to the nebular and stellar properties, particularly the mass of the progenitor star (e.g. Peimbert \\& Torres-Peimbert 1983, Calvet \\& Peimbert 1983, Stanghellini, Corradi \\& Schwarz 1993, Corradi \\& Schwarz 1995). This point of view is corroborated by the studies of Zhang \\& Kwok (1998). These authors, assuming a three-dimensional ellipsoidal shell with density variations and a complete absorption of the ionizing photons, which are reconverted in H recombination photons, produce simulated radio and H$\\alpha$ images which are compared to the observations. The visual appearence of PNe were also compared with models obtained for ionization-bounded prolate shells by Masson (1990). Assuming an ellipsoidal shell with constant thickness and that the ionizing radiation is geometrically diluted, Masson (1990) compared theoretical and observed H$\\alpha$ images. However, a given emission line image results from the sum of line intensities along the line of sight. The exact location of the regions which produce each line depends on the gas density distribution, but also on the ionization distribution, which, in turn, depends on the characteristics of the ionizing star and on the nebular gas abundance. Consequently, different emission lines provides information on different regions of the nebula, and only a consistent photoionization model applied to a realistic gas structure can give an accurate density distribution on the nebula.  In this paper we show that the geometry of a PN derived by fitting the observed image of a line emission and/or a line intensity ratio is misleading, since different geometries can reproduce the same observational data. Other constraints, coupled with a detailed 3D photoionization model, must be used for better defining the geometry, as illustrated by the planetary nebula NGC~3132.  Based on emission line imaging, an ellipsoidal geometry for NGC~3132 is assumed by various authors. Such a geometry is then suggested in a first approach. Studies similar to that of Masson (1990) using a prolate shell could explain the observed morphology of NGC~3132. Notice that the ellipsoidal shell model of Zhang \\& Kwok (1998) also reproduces the general trends of the observed H$\\alpha$ image of NGC~3132. Using a geometry defined by two concentric ellipsoids and a 3D photoionization code, B\\\"assgen, Diesch \\& Grewing (1990, hereafter \\cite{B90}) fit the observed [\\ion{O}{3}] and [\\ion{N}{2}] line images, as well as the emission line spectra taken through four slits.  Spatio-kinematic models for NGC~3132 were proposed by Sahu \\& Desai (1986, hereafter \\cite{SD86}) to explain the observed asymmetric double-peaked \\forb{O}{3}{5007} line profiles in five positions. \\cite{SD86} assert that an ellipsoidal model with velocity and density asymmetries reproduces the observed expansion velocities and correctly explains the nature of the profile asymmetries. However, in their paper, only the fit to the line profile at the central position is shown. Line profiles, as a powerful tool for studying the nebula geometry, are discussed by Morisset, Gruenwald \\& Viegas (2000, hereafter MGV2000).  As it will be shown below, all the models assuming an ellipsoidal emitting shell lead to the same problems: neither can reproduce the low central density observed by \\cite{SD86} and \\cite{J88} nor the observed asymmetric double-peaked profiles in regions far from the center of the nebula. After a description of the available observations of NGC~3132 (\\S \\ref{sec-obs}), we show that ellipsoidal shell models do not reproduce all the available observations (\\S \\ref{sec-bass}). A ``\\diab/'' nebula successfully explaining  all the main observational trends is presented in \\S \\ref{sec-diab}. The conclusions are outlined in \\S \\ref{sec-concl}. ",
        "conclusions": "\\label{sec-concl}  We have shown in this paper that the lack of high resolution observations (spatial and spectroscopic) may lead to a wrong conclusion about the morphology of a PN.  A new modeling tool based on a 3D photoionization code is used to consistently reproduce all the available observations of NGC~3132. Using an ellipsoidal geometry, all the low resolution observations (flux, emission line spectrum, imaging) could be reproduced. However, neither the density variations indicated by the \\dforb{S}{2}{6717}{6731} ratio, nor the \\forb{O}{3}{5007} velocity profiles could be explained by such a model. To account for these observations, a drastic change in the geometry of the nebula is necessary. The proposed Diabolo shape offers the solution to the density distribution, as well as to the asymmetric double-peaked emission line profiles, while also explaining the low resolution observations.  A Diabolo shape may reproduce the observed bipolar morphology of many PNe. In fact, Bryce, Balick \\& Meaburn (1994) suggested that NGC~6720 is not a spherical expanding shell as often assumed, but must be bipolar in nature, since the observed [NII] emission line profiles are split in two components. A bipolar shape emerges from studies of the formation and evolution of PNe, in particular those related to the interacting wind model (Kwok, Purton, \\& Fitzgerald 1978). Recent hydrodynamical models (e.g. Dwarkadas, Chevalier \\& Blondin 1996; \\cite{M97}) suggest that a Diabolo type shape can be produced in the accepted scenario of PNe formation (post AGB envelope ejection with wind interaction). In fact, recent high resolution images taken with the Hubble Space Telescope show that PNe with this type of structure is not uncommon. Some nice examples are the Hour-Glass nebula, NGC~2346, NGC~6537 and Hb-5.  No attempt was made to fine tuning the emission line ratios. This can be achieved through changes in the ionizing star characteristics, in the chemical abundances, as well as through small adjustments in the density distribution, which will not change the main conclusions of this paper, in particular the Diabolo shape. Furthermore, in the models presented here, the ionizing source is assumed to be at the center of symmetry of the nebula. However, there are evidences that the star should be shifted by $\\sim$ 1.7'' from the geometrical center of the nebula (Kohoutek \\& Lausten 1977 and \\cite{SD86}). A model accounting for this out-of-center ionizing star may still improve the results, leading to a better agreement with the observations regarding the brightness asymmetry and the emission line profile at the central position.  We have shown here the importance of a 3D photoionization model, combined with  imaging results and also with high resolution observations, to impose relevant constraints on the PN geometry. The misclassification of NGC~3132 as an ellipsoidal PN instead of a bipolar one casts doubts on the statistical results correlating the morphology and other PNe properties, such as abundances or progenitor type."
    },
    "0002/astro-ph0002140_arXiv.txt": {
        "abstract": "We present 21-cm \\HI\\ line and 13-cm continuum observations, obtained with the Australian Long Baseline Array, of the Seyfert 2 galaxy IC~5063.  This object appears to be one of the best examples of Seyfert galaxies where shocks produced by the radio plasma jet influence both the radio as well as the near-infrared emission.  The picture resulting from the new observations of IC~5063 confirms and completes the one derived from previous Australia Telescope Compact Array (ATCA) lower resolution observations.  A strong interaction between the radio plasma ejected from the nucleus and a molecular cloud of the ISM is occurring at the position of the western hot spot, about 0.6~kpc from the active nucleus.  Because of this interaction, the gas is swept up forming, around the radio lobe, a cocoon-like structure where the gas is moving at high speed.  Due to this, part of the molecular gas is dissociated and becomes neutral or even ionised if the UV continuum produced by the shocks is hard and powerful enough.  In the 21-cm \\HI\\ line new data, we detect only part of the strong blue-shifted \\HI\\ absorption that was previously observed with the ATCA at lower resolution.  In particular, the main component detected in the VLBI absorption profile corresponds to the most blue-shifted component in the ATCA data, with a central velocity of 2786 \\kms\\ and therefore blue-shifted $\\sim$614 \\kms\\ with respect to the systemic velocity.  Its peak optical depth is 5.4\\%.  The corresponding column density of the detected absorption, for a spin temperature of 100 K, is $N_{\\rm HI} \\sim 2 \\times 10^{21} $atoms cm$^{-2}$.  Most of the remaining blue-shifted components detected in the ATCA \\HI\\ absorption profile are now undetected, presumably because this absorption occurs against continuum emission that is resolved out in these high-resolution observations.  The \\HI\\ absorption properties observed in IC 5063 appear different from those observed in other Seyfert galaxies, where the \\HI\\ absorption detected is attributed to undisturbed foreground gas associated  with the large-scale  galaxy disk. In the case of IC~5063, only a small fraction of the absorption can perhaps be due to this.  The reason for this can be that the western jet in IC~5063 passes through a particularly rich ISM. Alternatively, because of the relatively strong radio flux produced by this strong interaction, and the high spectral dynamic range of our observations, broad absorption lines of low optical depth as detected in IC 5063 may have remained undetected in other Seyferts that are typically much weaker radio emitters or for which existing data is of poorer quality. ",
        "introduction": "The study of the effects of interactions between the radio plasma ejected from an active nucleus and the interstellar medium (ISM) of the hosting galaxy is presently attracting a lot of interest.  In particular, Seyfert and high-redshift radio galaxies appear to be the kind of objects where the effects of such interactions can be very important.  They can range from shaping the morphology of the gas in the ISM (with the radio plasma sweeping up material as it advances in the ISM), to the ionisation of the gas itself. While there is little doubt on the presence of such interactions in objects like Seyferts or high-$z$ radio galaxies, the actual importance of these effects in determining the overall characteristics of these sources is still a matter of debate.  In some Seyfert galaxies the morphological association between the radio plasma and the optical line-emitting clouds, as well as the presence of disturbed kinematics in these clouds, is striking.  In particular, the narrow-line regions (NLR) in Seyfert galaxies (i.e.\\ regions of highly ionised, kinematically complicated gas emission that occupy the central area -- up to $\\sim 1$ kpc from the nucleus) often appears to form a `cocoon' around the radio continuum emission (see e.g.\\ Wilson 1997 for a review; Capetti et al.\\ 1996; Falcke, Wilson \\& Simpson (1998) and references therein). Moreover, outflow phenomena are observed in the warm gas of several Seyfert galaxies (see Aoki et al.\\ 1996 for a summary).  Thus, the NLRs represent some of the best examples of regions where interaction between the local ISM and the radio plasma takes place and can be studied in detail.  The situation appears to be different for the atomic hydrogen.  Observations of the \\HI\\ 21-cm line, in absorption, can trace the distribution of this gas in front of the brightest radio components, that are usually observed in the central region of Seyferts (of kpc or sub-kpc size, i.e.\\ {\\sl co-spatial with the NLRs}).  Thus, the study of the distribution and kinematics of the {\\sl cold} component of the circumnuclear ISM can nicely complement the optical data.  Although \\HI\\ absorption has been detected in a number of Seyfert galaxies (e.g.\\ NGC~4151 Pedlar et al.\\ 1992; NGC~5929, Cole et al.\\ 1998; Mkn~6, Gallimore et al.\\ 1998 ; see also Brinks \\& Mundell 1996 and Gallimore et al.\\ 1999 and references therein), most of the investigated objects show single localised \\HI\\ absorption components that can be explained as rotating, inclined disks or rings aligned with the outer galaxy disk (Gallimore et al.\\ 1999) and only very seldom with gas in a parsec-scale circumnuclear torus (NGC~4151, Mundell et al.\\ 1995).  These components are therefore originated by gas that is not in interaction with the radio plasma.  However, more complex {\\sc H\\,i} absorption profiles that cannot be explained by the above mechanism have been observed in at least one Seyfert galaxy, IC~5063.  Australia Telescope Compact Array (ATCA) observations of this galaxy (Morganti, Oosterloo \\& Tsvetanov 1998, hereafter M98) have revealed a very interesting absorption system with velocities up to $\\sim 700$ \\kms\\ blue-shifted with respect to the systemic velocity.  In this object, unlike in other Seyfert galaxies, at least some of the observed \\HI\\ absorption is originating from regions of interaction between the radio plasma and the ISM, producing an outflow of the neutral gas. This object, therefore, poses a number of interesting questions as: where is the interaction occurring, what are the physical conditions, why such interaction is not  seen more often in neutral gas in other Seyfert galaxies?  Previous \\HI\\ observations were limited by low spatial resolution. In this paper  we present the results from new VLBI observations aimed at investigating in more detail its nuclear radio structure and locating where the complex \\HI\\ absorption observed with ATCA is really occurring.  Throughout the paper we adopt a Hubble constant of $H_\\circ = 50$ \\kms, so that 1~arcsec corresponds to 0.32 kpc at the redshift of IC~5063. ",
        "conclusions": "Using the Australian Long Baseline Array, we have detected a compact radio source of about of 50 mas (or $\\sim 16$ pc) in size (at 13 cm) in the Seyfert galaxy IC 5063.  Because of the high positional accuracy of these measurements, we can unambiguously identify this radio knot with the western radio lobe.  The hot spot is extended in a direction almost perpendicular to the radio jet.  In 21-cm line observations, we detect absorption very much blue-shifted ($\\sim$700 km s$^{-1}$) with respect to the systemic velocity.  Together with the 13-cm observations, this confirms that the \\HI\\ absorption is not taking place against the core, but that it is against the western radio knot.  At the position of the western radio knot a very strong interaction must be occurring between the radio jet and the ISM.  Various arguments suggest that this interaction is particularly strong compared to other Seyfert galaxies.  This makes IC 5063 a good candidate for studying the physics of jet-cloud interactions in Seyfert galaxies.  The HI absorption characteristics of IC 5063 are only partially consistent with other absorption studies of Seyfert galaxies. The major absorption component is occuring against the bright radio knot offset a few hundred parsecs from the core. While there are indication that the absorbing material is associated with the large scale \\HI\\ disk, it is clearly (and violently) disturbed by the passage of the jet. We suspect that more sensitive observations may reveal similar absorption profiles in other Seyfert galaxies with fainter radio sources."
    },
    "0002/astro-ph0002189_arXiv.txt": {
        "abstract": "Time-resolved spectroscopy of the novalike variable UU Aquarii is analyzed with eclipse mapping techniques to produce spatially resolved spectra of its accretion disc and gas stream as a function of distance from disc centre in the range 3600-6900 \\AA.  The spatially-resolved spectra show that the continuum emission becomes progressively fainter and redder for increasing disc radius -- reflecting the radial temperature gradient -- and reveal that the H\\,I and He\\,I lines appear as deep, narrow absorption features in the inner disc regions transitioning to emission with P Cyg profiles for intermediate and large disc radii.  The spectrum of the uneclipsed component has strong H\\,I and He\\,I emission lines plus a Balmer jump in emission and is explained as optically thin emission from a vertically extended disc chromosphere + wind. Most of the line emission probably arises from the wind. The spatially-resolved spectra also suggest the existence of gas stream ``disk-skimming'' overflow in UU~Aqr, which can be seen down to $R\\simeq 0.2\\; R_{L1}$. The comparison of our eclipse maps with those of Baptista, Steiner \\& Horne (1996) suggests that the asymmetric structure in the outer disc previously identified as the bright spot may be the signature of an elliptical disc similar to those possibly present in SU~UMa stars during superoutbursts. ",
        "introduction": "The standard picture of a novalike system is that of a close binary in which a late type star fills its Roche lobe and transfers matter to a companion white dwarf via an accretion disc. A bright spot is expected to form where the gas stream from the donor star hits the edge of the accretion disc.  The SW Sex stars (Thorstensen et~al. 1991) form a sub-class of the novalikes with orbital periods in the range 3-4 hs that do not seem to fit within the above standard picture, displaying a range of peculiarities: (1) single peaked asymmetric emission lines showing little eclipse, (2) large ($\\sim 70\\degr$) phase shifts between photometric and spectroscopic conjunction, (3) orbital phase-dependent absorption in the Balmer lines, (4) Doppler tomograms bright in the lower-left quadrant with small or no sign of disc emission, and (5) v-shaped continuum eclipses implying in flat radial temperature profiles in the inner disc (e.g., Warner 1995 and references therein). Earlier proposals to explain the phenomenon include accretion disc winds (Honeycutt, Schlegel \\& Kaitchuck 1986), magnetic white dwarfs disrupting the inner disc (Williams 1989), and gas stream overflow (Hellier \\& Robinson 1994). The two most recent models proposed to explain the phenomenology of the SW Sex stars are the disc-anchored magnetic propeller (Horne 1999) and a combination of stream overflow + disc winds (Hellier 1999).  UU~Aqr is an eclipsing novalike (P$_{\\rm orb}= 3.9$\\,hr) whose spectrum is dominated by single-peaked strong Balmer and He\\,I emission lines (e.g., Downes \\& Keyes 1988). H$\\alpha$ spectroscopy revealed that the line profile is highly asymmetric and phase-dependent and that the spectroscopic conjunction lags mid-eclipse by $\\sim 0.15$ cycle (Haefner 1989; Diaz \\& Steiner 1991). The lack of the rotational disturbance typical of emitting accretion discs during eclipse in H$\\beta$ led Hessman (1990) to the suggestion that the emission lines have a non-disc origin.  Baptista, Steiner \\& Cieslinski (1994; hereafter BSC94) derived a photometric model for the binary with $q= 0.30$, M$_1 = 0.67\\: {\\rm M}_{\\odot}$, an inclination of $i= 78$ degrees. From the analysis of mid-eclipse fluxes they suggested that the Balmer lines are formed in an extended region only partially occulted during eclipse, possibly in a wind emanating from the inner disc. They also found that UU Aqr presents long-term brightness variations of low amplitude ($\\simeq 0.3$ mags) on timescales of years.  The eclipse mapping study of Baptista, Steiner \\& Horne (1996; thereafter BSH96) indicates that the inner disc of UU Aqr is optically thick, resulting in a distance estimate of 200 pc. Temperatures in the disc range from $\\sim 6000$\\,K in the outer regions to $\\sim 16000$\\,K near the white dwarf at disc centre. The radial temperature profiles in the high state follow the T$\\:\\propto R^{-3/4}$ law in the outer and intermediate disc regions but flattens off in the inner disc, leading to mass accretion rates of $10^{-9.2}\\:{\\rm M_{\\odot}\\,yr}^{-1}$ at $R= 0.1\\: R_{\\rm L1}$ and $10^{-8.8}\\:{\\rm M_{\\odot}\\,yr}^{-1}$ at $R= 0.3\\: R_{\\rm L1}$ ($R_{\\rm L1}$ is the distance from disc centre to the inner Lagrangian point). Together with other characteristics, this led BSH96 to suggest that UU Aqr was an SW Sex star. The comparison of eclipse maps of the low and high states revealed that the differences are due to changes in the structure of the outer parts of the disc, the most noticeable effect being the appearance of a conspicuous red, bright structure at disc rim, which the authors identified with the bright spot.  According to BSH96, the \\.{M} of UU Aqr is barely above the critical limit for disc instability to set in. Warner (1997) noted that the outer disc temperature is only 6000 K and remarked that small variations in \\.{M} could lead to dwarf novae type outbursts. Honeycutt, Robertson \\& Turner (1998) performed a long-term photometric monitoring of UU Aqr which confirmed the high and low brightness states of BSC94 and revealed the existence of small amplitude ($\\simlt 1.0$ mag) brightness variations on timescales of a few days, which they called `stunted outbursts'.  The detailed spectroscopic study of Hoard et~al. (1998) reinforced the classification of UU Aqr as an SW Sex star. They found evidences for the presence of a bright spot at the impact site of the gas stream with the edge of the disc, and a non-axisymmetric, vertically and azimuthally extended absorbing structure in the disc. They proposed an explanation for the absorbing structure as well as for the other spectroscopic features of UU Aqr in terms of the explosive impact of the accretion stream with the disc. Optical and ultraviolet spectroscopy by Kaitchuck et~al. (1998) shows a secondary eclipse at phase 0.4 in the optical and Balmer lines (but not in the UV continuum or lines) which they suggested may be caused by an occultation of the bright spot and stream region by material suspended above the inner disc.  In this paper we report on the analysis of time-resolved spectroscopy of UU Aqr with multi-wavelength eclipse mapping techniques to derive spatially-resolved spectra of the accretion flow in this binary. Section~\\ref{data} describes the observations and data reduction procedures, while section~\\ref{analise} describes the analysis of the light curves with eclipse mapping techniques. Section~\\ref{resultados} presents eclipse maps at selected wavelengths, the radial intensity and brightness temperature distributions, spatially resolved spectra of the accretion disc and gas stream as well as the spectrum of the uneclipsed component. The results are discussed in section~\\ref{discussao} and summarized in section~\\ref{conclusao}. ",
        "conclusions": "\\label{conclusao}  We used time-resolved spectroscopy to study the structure and spectra of the accretion disc and gas stream of the novalike UU~Aquarii in the optical range. The main results of this analysis can be summarized as follows:  \\begin{itemize}  \\item The spectrum of the inner disc shows a blue continuum filled with deep, narrow absorption lines which transition to emission with clear P~Cygni profiles at intermediate and large radii ($R\\simgt 0.2 \\; R_{L1}$).  \\item The spectrum of the uneclipsed light has strong H\\,I and He\\,I emission lines and a Balmer jump in emission indicating a significant contribution from optically thin regions outside the orbital plane.  \\item Velocity-resolved eclipse maps and spectra indicate that most of the line emission probably arises in a vertically-extended disc chromosphere + wind.  \\item Differences in fractional contribution among emission lines suggests a vertical temperature gradient in the material above/below the disc.  \\item The comparison of the spectrum of the gas stream region and the disc region at the same radius as a function of radius gives evidence of gas stream disc-skimming overflow down to $R\\simeq 0.2\\; R_{L1}$. This may explain the phase-dependent absorption in emission lines.  \\item The comparison of our eclipse maps with those of BSH96 suggests that the asymmetric structure in the outer disc previously identified as the bright spot may be the signature of an elliptical precessing disc similar to those possibly present in SU UMa stars during superoutbursts.  \\end{itemize}"
    },
    "0002/astro-ph0002376_arXiv.txt": {
        "abstract": "Within the framework of the cold dark matter (CDM) cosmogony, a central cusp in the density profiles of virialized dark haloes is predicted. This prediction disagrees with the soft inner halo mass distribution inferred from observations of dwarf and low surface brightness galaxies, and some clusters of galaxies. By analysing data for some of these objects, we find that the halo central density is nearly independent of the mass from galactic to galaxy cluster scales with an average value of around $0.02 \\ M_{\\odot}/pc^3$. We show that soft cores can be produced in the CDM haloes by introducing a lower cut-off in the power spectra of fluctuations and assuming high orbital thermal energies during halo formation. However, the scale invariance of the halo central density is not reproduced in these cases. The introduction of self-interaction in the CDM particles offers the most attractive alternative to the core problem. We propose gravothermal expansion as a possible mechanism to produce soft cores in the CDM haloes with self-interacting particles. A global thermodynamical equilibrium can explain the central density scale invariance. We find  a minimum cross section capable of establishing isothermal cores in agreement with the observed shallow cores. If $\\sigma$ is the cross section, $m_x$ is the mass of the dark matter particle and $v$ is the halo velocity dispersion, then $\\sigma /m_x \\approx 4 \\ 10^{-25} (100 \\ km s^{-1}/v)$ $cm^2/GeV$. ",
        "introduction": "Constraining the nature of dark matter is presently one of the most relevant problems in cosmology and particle physics. The current most popular scenarios for structure formation in the universe are based on the inflationary CDM theory, according to which cosmic structures arise from small Gaussian density fluctuations composed mostly of non-relativistic collisionless particles. Luminous galaxies are thought to form by gas cooling and condensing into the dark matter haloes which grow by gravitational accretion and merging in a hierarchical fashion.  The question on the inner density profiles of the virialized dark matter haloes is at present controversial. In the last few years much observational and theoretical effort has been employed into investigating the inner structure of dark haloes. On galactic scales, the rotation curves of dwarf galaxies offer a way to study the inner mass distribution of their dark haloes directly since these galaxies are dominated by dark matter. By analysing the rotation curves of some near dwarf galaxies, Moore (1994), Flores $\\&$ Primack (1994) and  Burkert (1995) have shown that the central mass distribution of their dark haloes is soft, i.e. the haloes have a constant density core. A similar result concerns low surface brightness galaxies (LSB, hereafter) (de Blok $\\&$ McGaugh 1997) even though the uncertainty in the observational data is larger than in the case of dwarf spirals. Hern\\'{a}ndez $\\&$ Gilmore (1998) showed that the observed rotation curves of both LSB and normal large galaxies are consistent with a fixed initial halo shape, characterized by a significant soft core inner region. On scales of clusters of galaxies, unfortunately there is not much information available. Recently, from strong gravitational lensing observations, Tyson, Kochanski, \\& Dell'Antonio (1998) have obtained an unprecedent high-resolution mass map for the cluster CL0024+1654, which has not a central cD galaxy, and found the existence of a soft core. Taken together these studies suggest the univesality of constant density cores across both large mass scales and galactic types.  On the theoretical side, the structure of the CDM haloes was studied over a wide range of masses by means of high-resolution N-body cosmological simulations (e.g., Navarro, Frenk, \\& White 1997; NFW hereafter) and semi-analytical approaches (e.g., Avila-Reese, Firmani, \\& Hern\\'{a}ndez 1998). It was found that the universal density profile firstly introduced by NFW describes very well the mass distribution of most of the CDM haloes. This profile is univocally determined by the mass, and in the centre diverges as $\\rho \\propto r^{-1}$ producing a cusp in the core. Recent high-resolution N-body simulations have shown that, as the numerical resolution is increased, the inner profiles result even steeper than $r^{-1}$ (e.g., Moore et al. 1999b), making the CDM haloes more cuspy than in the case of the NFW profile.  So far, the predicted inner density profile of the CDM haloes seems to be in conflict with the observations. Another potential difficulty for the CDM models was recently reported: the N-body simulations predict an overly large number of haloes within group-like systems compared to observations (Klypin et al. 1999; Moore et al. 1999a). In light of these difficulties, the current stance of the hierarchical CDM-based scenario of structure formation remains somewhat confusing because, in fact, this scenario successfully accounts for: the distribution of matter at large scales (Bahcall et al. 1999), the uniformity of the cosmic microwave radiation and its small temperature anisotropies, and the observationally inferred cosmological parameters.  The aim of this letter is to analyse the halo core properties inferred from observations which might suggest explanations of the origin as to the soft halo cores and clarify the discrepancies that appear on small scales with the hierarchical scenario of structure formation. We investigate whether some modifications on the initial conditions of this scenario are able to improve the results with respect to the observations. We demonstrate that the introduction of self-interaction in the CDM particles as was suggested by Spergel \\& Steinhardt (1999) offers the most viable solution to the core problem in a context that preserves the hierarchical CDM-based scenario. ",
        "conclusions": "The discovery of a soft core in the cluster of galaxies CL0024+1254 by strong gravitational lensing measurements and the rotation curves of dark-matter dominated dwarf and LSB galaxies indicate that dark matter haloes have shallow inner density profiles from galactic to cluster scales. Studying in detail the observational data available for these cosmic objects, we found that the halo central density is nearly invariant with respect to the mass from galactic to cluster sizes.  We investigated different mechanisms and models for halo core formation within the hierarchical CDM scenario. We have shown that a lower cut-off at some wavelength in the CDM power spectrum and the assumption of high particle orbital thermal energies produce soft cores in the haloes, but the invariance of $\\rho_c$ with respect to the mass is not reproduced. A more viable solution to the core problem is the introduction of self-interaction in the CDM particles. Being this the case, we proposed the gravothermal expansion as the mechanism responsible for the formation of soft cores in a hierarchical CDM scenario.  Using a thermodynamical approach we have estimated the central density of haloes in the case of maximum efficiency for self- interaction and found good agreement with the values inferred from observations. The central density in this case scales with the halo mass and its maximum circular velocity as $\\rho _c\\propto V_m^6/M^2$. This result implies that $\\rho _c$ is roughly constant because for the CDM haloes $M\\propto V_m^n$ with $n\\approx 3.2$. If thermal equilibrium is restricted to the core, then the cross section given by eq.(2) may be derived consistently with observations. The cases analysed here, corresponding to a global and a local thermal equilibrium respectively, represent two limiting cases between which dark matter self-interaction may generate isothermal cores compatible with observations. We exclude from our analysis the extreme case of a {\\it very} strong self-interaction which may lead the core to a gravothermal catastrophe with a central density profile steeper than NFW. Such extreme assumption of large cross section may be immediately ruled out because a singular isothermal core will be produced in contradiction with observations.  We stress the relevance confirming the existence of soft cores with scale invariant densities would have. In particular, the construction of high-resolution mass maps with gravitational lensing techniques for the inner regions of clusters is of great interest."
    },
    "0002/astro-ph0002006_arXiv.txt": {
        "abstract": "The origin of the highest-energy cosmic rays remains a mystery. The lack of a high energy cutoff in the cosmic ray spectrum together with an apparently isotropic distribution of arrival directions have strongly constrained most models proposed  for the generation of these particles. An overview of the present state of theoretical proposals is presented. Astrophysical accelerators as well as top-down scenarios are reviewed along with their   most general signatures.   The origin and nature of these ultra-high energy particles will be tested by future observations and may indicate as well as constrain physics beyond the standard model of particle physics. ",
        "introduction": "The  detection of cosmic rays with energies above $10^{20}$ eV  has triggered considerable interest on  the  origin and nature of these particles. As reviewed by Watson \\cite{watson99} in this volume, many hundreds of events with energies above $10^{19}$ eV and about 20 events above $10^{20}$ eV  have now been observed by a number of experiments such as AGASA \\cite{takeda98,takeda99,Haya94}, Fly's Eye \\cite{bird},  Haverah Park \\cite{L91}, Yakutsk \\cite{Yak90}, and most recently the High Resolution Fly's Eye \\cite{hires}.  Most unexpected is the significant flux of events observed above $\\sim 7 \\times 10^{19}$ eV \\cite{takeda98}  with no sign of the Greisen-Zatsepin-Kuzmin (GZK) cutoff \\cite{GZK66}. A cutoff should be present if the  ultra-high energy particles are protons, nuclei, or photons from  extragalactic sources. Cosmic ray protons of energies above  a few $10^{19}$ eV lose energy to photopion production off the cosmic microwave background (CMB) and cannot originate further than about $50\\,$Mpc away from Earth. Nuclei are photodisintegrated on shorter distances due to the infrared background \\cite{PSB76SS99} while the radio background constrains photons to originate from even closer systems \\cite{bere70PB96}.  In addition to the presence of events past the GZK cutoff, there has been no clear counterparts identified in the  arrival direction of the highest energy events. If these events are protons,  cosmic ray observations should finally become astronomy!   At these high energies the Galactic and extragalactic magnetic fields do not affect their orbits significantly so that they should point back to their sources within a few degrees. Protons at $10^{20}$ eV propagate mainly in straight lines as they traverse the Galaxy since their gyroradii are $\\sim $ 100 kpc in $ \\mu$G  fields which is typical in the Galactic disk. Extragalactic fields are expected to be $\\ll \\mu$G \\cite{KronVallee,BBO99}, and induce at most $\\sim$ 1$^o$ deviation from the source. Even if the Local Supercluster has relatively strong fields, the highest energy events may deviate at most $\\sim$ 10$^o$ \\cite{RKB98,SLB99}.  At present, no correlations between arrival directions and plausible optical counterparts  such as sources in the Galactic plane, the Local Group, or the Local Supercluster have been clearly identified.  Ultra high energy cosmic ray (UHECR) data are consistent with an isotropic distribution of sources in sharp contrast to the anisotropic distribution of light within 50 Mpc from Earth.  The absence of a GZK cutoff and the isotropy of arrival directions are two of the many challenges that models for the origin of UHECRs face. This is an exciting open field, with many scenarios being proposed but no clear front runner.  Not only the origin of these particles  may be due to physics beyond the standard model of particle physics, but their existence  can be used to constrain extensions of the standard model such as violations of Lorentz invariance (see, e.g., \\cite{ABGG00}).  In the next section, a brief summary of the challenges faced by all theoretical models is given. In \\S3, astrophysical accelerators or ``bottom-up'' scenarios are reviewed, hybrid models are discussed in \\S4, and top-down scenarios  in \\S5.  To conclude, future observational tests of UHECR models and their implications are discussed  in \\S6. For previous reviews of UHECR models, the reader is encouraged to consult \\cite{hillas84,BBDGP90,gaisser90,bland99,berez99,BS99}. ",
        "conclusions": "Next generation experiments  such as the High Resolution Fly's Eye \\cite{hires2} which recently started operating, the Pierre Auger Project \\cite{cronin92} which is now under construction,  the proposed  Telescope Array \\cite{Teshima92}, and the OWL-Airwatch satellite \\cite{St97} will significantly improve the data at the extremely-high end of the cosmic ray spectrum \\cite{watson99}. With these observatories a clear determination of the spectrum and spatial distribution of UHECR sources is within reach. The lack of a GZK cutoff should become apparent with Auger \\cite{MT99a} and most extragalactic Zevatrons may be ruled out. The observed spectrum will distinguish Zevatrons from top-down models by testing power laws versus QCD fragmentation fits. The cosmography of sources should also become clear and able to discriminate  between plausible populations for UHECR sources. The  correlation of arrival directions  for events with energies above $10^{20}$ eV  with some known structure such as the Galaxy, the Galactic halo, the Local Group or the Local Supercluster would be key in differentiating between different models. For instance, a correlation with the Galactic center and  disk should become apparent at extremely high energies for the case of young neutron star winds \\cite{SH99}, while a correlation with the large scale galaxy distribution should become clear for the case of quasar remnants. If SHRs or monopolia are responsible for UHECR production, the arrival directions should correlate with the dark matter distribution and show the halo asymmetry. For these signatures to be tested, full sky coverage is essential. Finally,  an excellent discriminator would be an unambiguous composition determination  of the primaries. In general, Galactic disk models  invoke iron nuclei to be consistent with the isotropic distribution,  extragalactic Zevatrons tend to favor proton primaries, while photon primaries are more common for early universe relics. The hybrid detector of the Auger Project should help determine the composition by measuring simultaneously the depth of shower maximum and the muon content of the same shower.   In addition to explaining the origin of UHECRs, GUT to Planck scale physics can potentially be probed by the existence of UHECRs. For instance, the breaking of Lorentz invariance can change the threshold for photopion production significantly in such a way as to be constrained by a clear observation of the GZK cutoff \\cite{ABGG00}.   There are great gains to be made if the data at the highest energies is improved by a few orders of magnitude. The prospect of testing extremely high energy physics as well as solving the UHECR puzzle given all the presently proposed  models sends a strong message that the challenge is back in the observational arena. Fortunately, observers have accepted the challenge and are building and planning  experiments  large enough to resolve  these open questions \\cite{watson99}."
    },
    "0002/astro-ph0002283_arXiv.txt": {
        "abstract": "FIRBACK is a one of the deepest surveys performed at 170 $\\mu m$ with ISOPHOT onboard ISO, and is aimed at the study of cosmic far infrared background sources. About 300 galaxies are detected in an area of four square degrees, and source counts present a strong slope of 2.2 on an integral \"logN-logS\" plot, which cannot be due to cosmological evolution if no K-correction is present. The resolved sources account for less than 10\\% of the Cosmic Infrared Background at 170 $\\mu m$. In order to understand the nature of the sources contributing to the CIB, and to explain deep source counts at other wavelengths, we have developed a phenomenological model, which constrains in a simple way the luminosity function evolution with redshift, and fits all the existing deep source counts from the mid-infrared to the submillimetre range. ",
        "introduction": "The Cosmic Infrared Background (CIB), due to the accumulation of galaxy emission at all redshifts along the line of sight in an instrument beam, is a powerful tool for studying galaxy evolution. FIRBACK (\\cite{puget99a}, \\cite{dole99}), one of the deepest surveys performed at $170\\,\\mu m$, is aimed at the study of the CIB, in two complementary ways: \\begin{itemize} \\vspace{-0.2cm} \\item study the resolved sources (this paper) \\item study the background fluctuations (\\cite{lagache99a}, \\cite{lagache99b}) \\end{itemize} \\vspace{-0.2cm} Throughout this paper we use a cosmology with $h=0.65$, $\\Omega=1$ and $\\Lambda=0$. ",
        "conclusions": "We have presented results from the FIRBACK survey, one of the largest ISO programs, dealing with resolved sources of the CIB: our source counts at $170\\,\\mu m$ show strong evolution. This evolution is explained by a simple law involving the redshift of the luminosity function, which differs from pure density or luminosity evolution: the ``Bright End Luminosity Function Evolution'' model. The model fits all the existing source counts at 15, 90, 170 and 850 $\\mu m$, and also predicts a redshift distribution in agreement with the (sometimes sparse) observations. This powerful tool, based on observational constraints on the CIB spectrum inversion and the local Luminosity Function, and on assumed template galaxy spectra, not only agrees with existing data but also is able to make useful predictions on source counts or CIB fluctuations.  These predictions may be useful for planning the utilization of major telescopes of the future, such as SIRTF, Planck, FIRST and ALMA. All the FIRBACK materials are available at: \\verb|http://wwwfirback.ias.fr|. \\vspace{-0.3cm}"
    },
    "0002/hep-ph0002063_arXiv.txt": {
        "abstract": "We point out that for solar neutrino oscillations with the mass--squared difference of $\\Delta m^2 \\sim 10^{-10}- 10^{-9}$ eV$^2$, traditionally known as ``vacuum oscillation'' range, the solar matter effects are non-negligible, particularly for the low energy \\textsl{pp} neutrinos. One consequence of this is that the values of the mixing angle $\\theta$ and $\\pi/2-\\theta$ are not equivalent, leading to the need to consider the entire physical range of the mixing angle $0\\le\\theta\\le\\pi/2$ when determining the allowed values of the neutrino oscillation parameters. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0002/hep-ph0002255_arXiv.txt": {
        "abstract": "This is somewhat extended version of the talk given at the Gran Sasso Summer Institute: Massive Neutrinos in Physics and Astrophysics. It describes general ideas about mirror world, extra spatial dimensions and dinosaur extinction. Some suggestions are made how these seemingly different things can be related to each other. ",
        "introduction": "The history of science in particular and the Human history in general teach us it is not easy to answer a simple question \\cite{1} ``What is truth?''. Maybe because the truth usually has infinitely many aspects or projections to be grasped. The following example from \\cite{2} I like very much. Consider the two figures below. \\begin{figure}[htb] \\begin{center} \\mbox{\\epsfig{figure=tri.eps ,  height=4.0cm}} \\end{center} \\label{Fig1} \\end{figure}  \\noindent Are they different? Certainly they are. This seems to be an indisputable truth. But if our wonderful divine gift -- our imagination -- will help us to escape bonds of the stiff two-dimensional logic we can see the following three-dimensional picture: \\begin{figure}[htb] \\begin{center} \\mbox{\\epsfig{figure=con.eps ,  height=6.0cm}} \\end{center} \\label{Fig2} \\end{figure}  \\noindent Now it is clear (an indisputable truth again) that these two figures, which originally appeared as two different objects, actually are just two different projections of the same thing -- the cone. From the original two-dimensional picture alone it is impossible to establish with certainty whether these two figures are substantially different or not. With the help of imagination we can find as a viable option that these figures may have the common origin and represent in fact the one essence, but in order to prove the case one needs further information (some experiments?)  What follows is an attempt to answer Bludman's question \\cite{3} ``Mu\\ss{} es sein?'' with regard to the Mirror World. Until experiments firmly prove or disprove its existence, any answer will include a great deal of imagination by necessity. So I will describe things at first sight very different and not related to the Mirror World. I refer to your imagination to accept a possibility that these different tales are in fact fragments of the same story.  To demonstrate the importance of imagination, I will perform a little hocus-pocus now and find the Mirror World even in a simple arithmetical expression. ",
        "conclusions": "The mirror Nemesis hypothesis emerged almost as a joke during our e-mail discussions with Robert Foot. After some thought we found no reason why this hypothesis, although somewhat extravagant, might not be true \\cite{80}. As a result, the dinosaur theme, which I originally intended to introduce for just to make presentation more vivid, quickly became one of the central motives of this talk, and you have the story presented above. I hope you enjoyed it regardless whether dinosaurs were really eyewitnesses of the mirror world or not.  I consider the possibility to restore the equivalence between left and right through the mirror world as very attractive. Theories with extra spatial dimensions, and $M$-theory in particular, can easily produce various ``shadow worlds'' which are however not necessarily parity invariant (this refers to the $E_8\\times E_8$ model also, mentioned earlier), but some of them might be, so realizing the mirror world scenario. Maybe, it is even possible to have a mirror world without mirror particles. $M$-theory nonorientable compactifications, suggested so far \\cite{81}, do not lead to the realistic model, as I can judge. But it will be very interesting to find a realistic example and show that the parity noninvariance of our world indicates to its nonorientable topology and is only local phenomenon.  ``In the Soviet scientific society the scientists had one freedom that scientists in the West lacked and still lack (perhaps the only real freedom that Eastern scientists had), and that was to spend time also on esoteric questions. They did not have to be scrutinized by funding agencies every now and then'' \\cite{82}. The Soviet Union disappeared and so did this freedom. You can consider this paper, if you like, as a nostalgia for this kind of freedom, enabling to escape bonds of the stiff pragmatic logic."
    }
}